The Luminosity Distribution of Globular Clusters in Dwarf Galaxies

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The mRNA-Bound Proteome and Its Global Occupancy Profile on Protein-Coding Transcripts

Molecular CellResourceThe mRNA-Bound Proteome and Its Global Occupancy Proﬁle on Protein-Coding TranscriptsAlexander G.Baltz,1,3Mathias Munschauer,1,3Bjo¨rn Schwanha¨usser,1Alexandra Vasile,1Yasuhiro Murakawa,1 Markus Schueler,1Noah Youngs,2Duncan Penfold-Brown,2Kevin Drew,2Miha Milek,1Emanuel Wyler,1Richard Bonneau,2Matthias Selbach,1Christoph Dieterich,1and Markus Landthaler1,*1Max Delbru¨ck Center for Molecular Medicine,Berlin Institute for Medical Systems Biology,13125Berlin,Germany2Center for Genomics and Systems Biology,Department of Biology,New York University,New York,NY10003,USA3These authors contributed equally to this work*Correspondence:ndthaler@mdc-berlin.deDOI10.1016/j.molcel.2012.05.021SUMMARYProtein-RNA interactions are fundamental to core biological processes,such as mRNA splicing,locali-zation,degradation,and translation.We developed a photoreactive nucleotide-enhanced UV crosslink-ing and oligo(dT)puriﬁcation approach to identify the mRNA-bound proteome using quantitative pro-teomics and to display the protein occupancy on mRNA transcripts by next-generation sequencing. Application to a human embryonic kidney cell line identiﬁed close to800proteins.To our knowledge, nearly one-third were not previously annotated as RNA binding,and about15%were not predictable by computational methods to interact with RNA. Protein occupancy proﬁling provides a transcrip-tome-wide catalog of potential cis-regulatory regions on mammalian mRNAs and showed that large stretches in30UTRs can be contacted by the mRNA-bound proteome,with numerous putative binding sites in regions harboring disease-associ-ated nucleotide polymorphisms.Our observations indicate the presence of a large number of mRNA binders with diverse molecular functions partici-pating in combinatorial posttranscriptional gene-expression networks.INTRODUCTIONDuring and immediately after transcription,nascent messenger RNAs(mRNAs)associate with proteins to form messenger ribo-nucleoprotein(mRNP)complexes that mediate and regulate most aspects of mRNA metabolism and function.Throughout their life cycle,mRNP complexes consist of a dynamically changing repertoire of proteins that deﬁne processing,cellular localization,and the decay and translation rate of mRNAs.The mammalian genome has been predicted to encode about 600RNA-binding proteins(de Lima Morais et al.,2011),based on the presence of one or more catalytic or noncatalytic domains that can interact with RNA.However,several proteins implicated in other cellular processes exhibit RNA-binding activity despite the absence of recognizable RNA-binding domains.Among them,cytosolic aconitase(also known as iron-regulatory protein1)posttranscriptionally regulates speciﬁc target mRNAs depending on cellular iron levels(Kennedy et al.,1992).This and other examples of RNA-binding activity of unexpected proteins highlight the need to thoroughly catalog the cellular repertoire of RNA-binding proteins in order to deﬁne the system that regu-lates the posttranscriptional fate of mRNAs.More than30years ago,initial attempts were made to isolate and analyze the poly(A)+RNA-bound proteome by oligo(dT)sepharose chromatography.Puriﬁcations of mRNPs from in vitro UV-irradiated polysomal fractions(Greenberg, 1979),from UV-irradiated intact cells(Wagenmakers et al., 1980),and from untreated cells(Lindberg and Sundquist,1974) revealed the association of a speciﬁc set of proteins with ter on,similar methods were applied to characterize hnRNP particles and to identify the mRNA polyadenylate-binding protein(Adam et al.,1986;Choi and Dreyfuss,1984). Recently,screening approaches and oligo(dT)puriﬁcation procedures were used to provide a catalog of yeast RNA-binding proteins(Scherrer et al.,2010;Tsvetanova et al.,2010). However,methods for comprehensive identiﬁcation of mamma-lian mRNA-binding proteins have remained elusive.Another prerequisite for our understanding of the function of RNA-interacting proteins is a systematic identiﬁcation of their binding sites and the deﬁnition of their RNA targets.Current genomic approaches use UV crosslinking and immunoprecipi-tation(CLIP)of mRNA-protein complexes in combination with next-generation sequencing to identify RNA binding sites(Ko¨nig et al.,2010;Licatalosi et al.,2008).One recently developed method,PAR-CLIP,employs the photoreactive thionucleo-sides,4-thiouridine and6-thioguanosine,to increase the cross-linking efﬁciency between protein and RNA and to provide near nucleotide resolution of the RNA-binding site(Hafner et al., 2010).Here,we use a photoreactive nucleoside-enhanced UV cross-linking and oligo(dT)afﬁnity puriﬁcation approach to identify the mRNA-bound proteome and to globally map the sites of protein-mRNA interactions in mammalian ing high-resolution quantitative mass spectrometry we uncovered close to800 proteins in oligo(dT)-puriﬁed protein-mRNA complexes.A large number of these proteins were neither previously annotated674Molecular Cell46,674–690,June8,2012ª2012Elsevier Inc.nor could be predicted by state of the art computational methods to interact with RNA.We validated the RNA binding function of 19of these mRNA-interacting ing PAR-CLIP-seq, we identiﬁed mRNA binding preferences and distinct binding patterns forﬁve of these proteins.Furthermore,protein occu-pancy proﬁling on mRNA by next-generation sequencing of protein-crosslinked RNA fragments,provided a transcriptome-wide view of the interaction sites of the mRNA-bound proteome and revealed widespread binding of proteins to30untranslated regions(30UTRs)of mRNAs.RESULTSOptimization of mRNP Oligo(dT)Afﬁnity PuriﬁcationTo characterize the protein-mRNA interactome,we sought to improve existing methods to identify the protein content of oligo(dT)afﬁnity-puriﬁed mRNP complexes and to determine the mRNA regions contacted by the mRNA-bound proteome (Figure1A).A key feature of our approach is the use of photo-reactive nucleoside analogs,4-thiouridine(4SU)and6-thio-guanosine(6SG),to metabolically label cellular RNA without detectable incorporation into DNA(Figure S1A available online) (Favre et al.,1993).Both4SU and6SG are readily taken up by cultured mammalian cells and dramatically enhance the cross-linking efﬁciency of proteins to RNA by UV365nm irradiation (Hafner et al.,2010).Photocrosslinking of living cells stabilizes mRNP complexes and facilitates their isolation by oligo(dT) afﬁnity puriﬁcation(Setyono and Greenberg,1981;Wagen-makers et al.,1980).Protein-denaturing conditions during the puriﬁcation ensure a stringent isolation of proteins in direct contact with mRNA through covalent bonds and thus enable the identiﬁcation of the mRNA-interacting proteins by mass spectrometry(Figure1A).Moreover,4SU-labeled RNA,cross-linked to proteins,can readily be identiﬁed by characteristic T to C(T-C)transitions in complementary DNA(cDNA)(Hafner et al.,2010),providing a way to globally identify the RNA segments that are crosslinked by the mRNA-bound proteome (Figure1A).We initially tested this approach by purifying protein-mRNA complexes using oligo(dT)beads from extracts of UV-irradiated and nonirradiated intact human embryonic kidney(HEK)293 cells after growth in medium supplemented with or without 4SU and6SG.Resolving the RNase-treated eluate by SDS-PAGE revealed that the combination of metabolic labeling of RNA with photoreactive nucleosides and irradiation at UV 365nm allowed a high recovery of proteins(Figure1B).We further examined the amount of mRNA obtained in precipitates from extracts of crosslinked and noncrosslinked cells.A qRT-PCR analysis showed that comparable amounts of GAPDH mRNA were precipitated(Figure S1B),suggesting that labeling of RNA and UV irradiation had only a minor effect on the mRNA pull-down efﬁcacy.As expected when probing the oligo(dT)precipitate for the presence of known RNA-binding proteins by western analysis, we were able to detect the heterogeneous nuclear ribonucleo-protein K(HNRNPK)(Figure S1C).However,the Argonaute protein,AGO2/EIF2C2,was not detectable after a single oligo(dT)puriﬁcation(Figure S1C),likely due the insufﬁcient precipitation of mRNAs and/or incomplete capture of mRNAs with shortened poly(A)tails,such as microRNA/AGO-targeted mRNAs.Thus,we measured the degree of depletion of GAPDH mRNA after one oligo(dT)precipitation.The GAPDH transcript is abundant and targeted by AGO proteins(Hafner et al.,2010; Kishore et al.,2011).Figure1C shows that only about70%of this transcript was depleted in the supernatant when compared to input RNA.Three additional consecutive pull-downs from the same extract reduced the amount of GAPDH mRNA in the super-natant to about5%(Figure1C).A western analysis of the pooled eluates of four oligo(dT)puriﬁcations validated the presence of AGO2protein(Figure1D)as well as the RNA-binding protein QUAKING,indicating that multiple consecutive oligo(dT)precip-itations are required to precipitate crosslinked AGO protein. Known DNA-binding proteins were not detected in the pooled precipitates(Figure S1D).Characterization of the Oligo(dT)-Puriﬁed RNATo obtain a more detailed picture of the RNA present in the pooled precipitates of four consecutive oligo(dT)puriﬁcations, we constructed a cDNA library by random priming of4SU-and 6SG-labeled RNA derived from irradiated and nonirradiated cells.Digital gene expression analysis of the cDNA library of nonirradiated cells,labeled with4SU and6SG,revealed that about88%of the sequence reads mapped to mRNA and8% ribosomal RNA(rRNA)genes,whereas in RNA precipitates ob-tained from UV-irradiated cells the rRNA content increased to 36%,likely reﬂecting crosslinking of ribosomes to mRNA tran-scripts(Figures1E and S1E).In contrast a standard mRNA puri-ﬁcation procedure,involving only a single oligo(dT)precipitation, of untreated cells resulted in96%mRNA and2%rRNA(Figures 1E and S1E).Furthermore a comparison of the different RNA libraries showed that the abundance of mRNAs obtained by a single oli-go(dT)puriﬁcation from untreated cells and metabolically labeled transcripts derived from noncrosslinked and UV-crosslinked cells correlated well(Pearson correlation coefﬁcient of0.87 and0.82,respectively;Figure1F),indicating that the oligo(dT)-precipitated mRNA closely reﬂected the cellular mRNA pool. To monitor the incorporation of photoreactive nucleotides into mRNA,we isolated4SU-and6SG-labeled RNA from the oli-go(dT)precipitate of noncrosslinked cells by biotinylation and streptavidin puriﬁcation(Do¨lken et al.,2008).The abundance of the thionucleotide-containing RNA was in good agreement with cellular mRNA(Pearson correlation coefﬁcient of0.90),sug-gesting efﬁcient and unbiased metabolic labeling of transcripts (Figure1F).In summary,we concluded that at least four consec-utive oligo(dT)puriﬁcations are required to efﬁciently purify cellular mRNA-protein complexes,while a signiﬁcant enrichment of mRNA over other classes of RNA can still be observed in the resulting precipitates.Identiﬁcation of mRNA-Bound Proteins by Quantitative Mass SpectrometryTo identify proteins crosslinked to mRNAs,we performed oligo(dT)puriﬁcations,as described above,and precipitates were analyzed by SILAC-based quantitative mass spectrometry (Mann,2006).For this purpose,cells were grown in mediumMolecular CellThe Protein-mRNA InteractomeMolecular Cell46,674–690,June8,2012ª2012Elsevier Inc.675Figure 1.Analysis of Oligo(dT)-Puriﬁed mRNPs(A)Illustration of the experimental setup to identify the mRNA-bound proteome and its occupancy proﬁle on RNA.Transcripts were labeled with photoreactive nucleosides and proteins were crosslinked to RNA by 365nm UV irradiation.mRNP complexes were isolated after cell lysis by oligo(dT)precipitation under denaturing conditions.For the identiﬁcation of the mRNA-bound proteome,mRNPs were eluted from the beads,nuclease treated,and analyzed by quantitative mass spectrometry.For identiﬁcation of the protein binding pattern on RNA,mRNPs were RNase I treated,followed by proteinase K digest to remove RNA-bound proteins.RNA molecules were converted into a cDNA library and next-generation sequenced.(B)SDS-PAGE analysis of proteins crosslinked to polyadenylated RNA.HEK293cells were grown in medium supplemented with 4SU and/or 6SG and UV irradiated at 365nm.Cells were lysed using denaturing conditions,and protein-mRNA complexes were isolated by oligo(dT)precipitation.Protein-RNA complexes were eluted from oligo(dT)beads,treated with RNase I,separated on a SDS gradient gel,and visualized by silver staining.(C)GAPDH mRNA depletion.qRT-PCR analysis of GAPDH mRNA in supernatants (SN1to SN4)after each round of oligo(dT)bead precipitation (four in total)compared to GAPDH mRNA in extract before precipitations (input)are shown as percent of input.The error bars display the calculated maximum and minimum expression levels that represent the standard error of the mean expression level with a 95%conﬁdence interval.Molecular CellThe Protein-mRNA Interactome676Molecular Cell 46,674–690,June 8,2012ª2012Elsevier Inc.supplemented with ‘‘light’’or ‘‘heavy’’stable isotope labeled amino acids to compare the protein abundance in oligo(dT)precipitates of crosslinked cells to that of noncrosslinked cells (Figure S2A).We performed two independent experiments (L1and L2)in which the ‘‘light’’labeled cells were UV irradiated and proteins in the oligo(dT)precipitations were compared to the precipitate of noncrosslinked ‘‘heavy’’labeled cells.In a single ‘‘label swap’’experiment (H1),the ‘‘heavy’’labeled cells were crosslinked and the recovered proteins were compared to those of ‘‘light’’labeled noncrosslinked cells.Proteins in pooled oligo(dT)precipitates were separated by SDS-PAGE followed by in-gel digest.Peptides were analyzed by nanoﬂow high-performance liquid chromatography (HPLC)and online mass spectrometry on a high-resolution instrument.In total,300hr of data acquisition for all three experiments re-sulted in 1,383,856MS/MS spectra.Processing the data with MaxQuant (Cox and Mann,2008),we identiﬁed 1,326proteins and observed a signiﬁcant overlap between experiments.Seven hundred ninety proteins were identiﬁed in all three proteomic analyses,and 561of those were quantiﬁed with at least three observed SILAC-peptide ratios in each experiment (Figure 2A).To further examine the reproducibility we compared log2SILAC ratios from biological replicates L1and L2(Figure 2B).Seven hundred seventy-eight out of 800proteins identiﬁed in both experiments were speciﬁcally enriched in the precipitates of UV-crosslinked cells relative to the nonirradiated control cells (SILAC log2fold changes <0in both cases).Hence,97%of all identiﬁed proteins showed speciﬁc enrichment.In addition we observed no correlation between the fold enrichment and the cellular protein abundance (Figure S2B),suggesting that the degree of enrichment was independent of the number of protein molecules present in the cell.Next,we plotted the log2SILAC ratios from both biological replicates against the label swap experiment.As expected,most SILAC ratios were inverted by the label swap (Figures 2C and 2D).The proteins with low SILAC ratios in both the biological replicates and the label swap experiment were largely contami-nants such as trypsin,LysC,and keratins.We therefore excluded 135proteins with negative log2SILAC ratios in the label swap experiment.Among the excluded proteins were six known RNA-binding proteins:the small SNRPE,PDCD11,ELAVL3,RBM16,PA2G4,and RBPMS.Requiring an enrichment of at least 3-fold (determined by SILAC ratios)in at least one of three analyses,we excluded 54proteins,which did not meet this strin-gent enrichment criteria,thereby ending up with 889proteins.In order to facilitate further analysis,we retained only the longestisoforms of the identiﬁed proteins,resulting in a nonredundant list of 797proteins (Table S1).We further subdivided the 797proteins into three classes (Table S1).Class I included 505proteins (63%),which were enriched more than 3-fold in all three proteomic analyses.One hundred ninety proteins (24%)showed a 3-fold enrichment in two experiments (class II),and 102proteins (13%)in only one experiment (class III).Overview of Identiﬁed mRNA-Interacting ProteinsWe ﬁrst classiﬁed the identiﬁed proteins into functional cate-gories based on gene annotation.As expected,ribosomal proteins,RNA helicases,translation factors,and RNA-binding proteins were most frequent,making up close to 70%of the identiﬁed proteins (Figure 3A).The low numbers of highly ex-pressed cellular proteins such as metabolic enzymes,histones,and heat-shock proteins,suggested that the oligo(dT)puriﬁca-tion was speciﬁc.The median relative abundance of RNA-binding proteins,translations factors,and ribosomal proteins was slightly higher but comparable to that of other functional groups of proteins (Figure 3B).Conﬁrming the method,we discovered RNA-interacting proteins present in complexes that inﬂuence surveillance and translation of spliced mRNAs.We detected all proteins,RBM8A/Y14,MAGOH,EIF4A3,and CASC3/BTZ,making up the core of the exon junction complex (Bono et al.,2006).However,we identiﬁed an order of magnitude more molecules of EIF4A3and CASC3(Table S1),which were previously shown to directly interact with RNA,compared to MAGOH and the RNA-recognition motif-containing protein RBM8A.We observed EIF4A1,EIF4B,EIF4E,EIF4G1,and EIF4H,all of which are present in the translation initiation complex (Jackson et al.,2010).Additionally,we identiﬁed three of the four Argonaute proteins expressed in HEK293cells.On the other hand,the iden-tiﬁed mRNA binders only partially overlapped with sets of proteins found in nuclear RNA-containing structures.Ninety-nine out of 172proteins detected in spliceosomal complexes (Bessonov et al.,2008)were observed to interact with mRNA (Figure 3C).Two hundred forty-one identiﬁed mRNA interactors were also found in the nucleolus proteome (Andersen et al.,2005)(Figure 3C).In addition to the expected mRNA-interacting proteins,we identiﬁed 245proteins (Table S1),which have not been previ-ously annotated as RNA binding (Figure 3A,candidate mRNA binders).Eighty percent of these proteins were detected in at least two out of three proteomic analyses,and about 47%(D)Western analysis of FLAG/HA-tagged RNA-binding proteins QUAKING (QKI)and ARGONAUTE 2(AGO2/EIF2C2)in input extract (I),supernatant after precipitation (S),and oligo(dT)-puriﬁed material (P)of UV-crosslinked and noncrosslinked cells.(E)Read count distribution over different RNA types.mRNA was puriﬁed either from total TRIzol-extracted RNA by a single oligo(dT)precipitation (mRNA seq),or by four rounds of oligo(dT)precipitation from cellular extract of UV-irradiated and nonirradiated cells (4SU+6SG UV and 4SU+6SG no UV,respectively).Crosslinked proteins were removed by Proteinase K digest prior to RNA analysis by next-generation sequencing of recovered RNA.The read count distribution over different RNA classes (mRNA,rRNA,and other)was inferred by multiplication of the FPKM values with the respective length of the longest transcript of a given gene.(F)Pair-wise correlation between RNA abundance expressed as log2FPKM of RNA described in (E).For assessment of the incorporation of photoreactive nucleoside into RNA,the 4SU-and 6SG-containing RNA was puriﬁed from oligo(dT)-precipitated RNA of noncrosslinked cells by biotinylation and streptavidin pull-down (4SU+6SG puriﬁed RNA)and analyzed by next-generation sequencing.The diagonal is shown as yellow line for of each pairwise comparison,whereas a LOESS regression line is shown in red.See also Figure S1.Molecular CellThe Protein-mRNA InteractomeMolecular Cell 46,674–690,June 8,2012ª2012Elsevier Inc.677were observed in all three pull-downs.We applied an adaptation of a multiple association network integration algorithm (Mosta-favi et al.,2008)to predict proteins with RNA-binding function,using gene ontology (GO)data,InterPro and Pfam domain data,gene coexpression,protein-protein interaction,and structural similarity data (Drew et al.,2011).This algorithm demonstrated strong predictive power,as evidenced by the precision-recall values for RNA binding (see Table S2,part A)and by previous ﬁeld-wide tests of function prediction algorithms(Pen˜a-Castillo et al.,2008).The most informative data types for predicting RNA-binding were determined to be GO process,GO localization,and 3D protein structure similarity,followed to a lesser degree by protein-protein interaction data,gene coex-pression,and InterPro domains.A full description of the algo-rithm and benchmarking results appear in the Supplemental Experimental Procedures .After applying the algorithm to the 245candidate mRNA-interacting proteins detected by our assay,112proteins(48belonging to class I)could not be predicted as RNA binders (Table S1)(even at a very low precision level R 20%,and when using the function prediction algorithm in a manner that mini-mizes false negative predictions at the expense of false positive predictions).This ﬁnding strongly suggests that our experiments uncovered RNA-interacting proteins that use distinct or highly divergent RNA-binding domains and occupy discrete regions of the known protein association networks.Some of our discov-eries include proteins that are functionally annotated as tran-scription factors (JUN,NXF1),and putative protein kinases (FASTKD1,FASTKD2,FASTKD5).Additionally,several pro-teins encoded by open reading frames (C1orf35,C16orf80,C11orf31,C9orf114,C19orf47)were observed to be RNA binding.Overrepresentation of Nucleic Acid Binding Domains Next,we classiﬁed the identiﬁed proteins based on their three-dimensional structure and amino acid sequence.FortheFigure 2.Identiﬁcation of the mRNA-Bound Proteome by Quantitative Mass Spectrometry(A)Summary of proteomic experiments.In two replicates,the proteomic composition of oligo(dT)precipitates was analyzed for ‘‘light’’labeled crosslinked cells (experiments L1and L2)and one experiment for ‘‘heavy’’labeled crosslinked cells (H1).The overlap of identiﬁed proteins in different experiments is shown in the Venn diagram.The table indicates the number of identiﬁed and quantiﬁed proteins,as determined by SILAC ratios of proteins in each experiment.(B)Comparison of the log2fold changes (LFC)of ‘‘heavy’’to ‘‘light’’SILAC ratios (H/L)of proteins quantiﬁed in biological replicates L1and L2.Previously known RNA-binding proteins are indicated in green,and known contaminants are in red.(C and D)As in (B),proteins were quantiﬁed in L1(C)or L2(D)plotted against proteins quantiﬁed in label swap experiment H1.See also Figure S2and Table S1.Molecular CellThe Protein-mRNA Interactome678Molecular Cell 46,674–690,June 8,2012ª2012Elsevier Inc.Figure 3.Overview of Identiﬁed mRNA-Interacting Proteins(A)Number of identiﬁed proteins belonging to different functional categories.(B)Median relative number of protein molecules belonging to different functional categories as determined,shown as box plots.Protein amounts were calculatedas the sum of all peptide peak intensities divided by the number of theoretically observable tryptic peptides (Schwanha¨usser et al.,2011).The median is shown as horizontal line,and the surrounding box deﬁnes the upper and lower quartile.The sample range is deﬁned by the whiskers,while dots indicate potential outliers.(C)Overlap of identiﬁed mRNA binders with proteins present in the spliceosome and the nucleolus.(D)Number of identiﬁed proteins with speciﬁc RNA-binding domains (dark gray)was compared to the respective number of RNA-binding domain-containing proteins in HEK293proteome (light gray).(E)mRNA-bound proteins and their ﬁrst neighbors based on PPI were analyzed for overrepresented Gene Ontology terms.The members of the cluster enriched for DNA damage response are depicted as nodes in the PPI network.Only proteins with direct interactions are shown.The node color represents the number of times the proteins were identiﬁed in our analysis (purple,detected in all three experiments;red,detected in two experiments;pink,detected in one experiment),and the edges indicate direct PPIs between the cluster members.The node shape indicates the level of prediction.Annotated RNA-binding proteins are marked by green node border color.See also Figure S3and Tables S1and S3.Molecular CellThe Protein-mRNA InteractomeMolecular Cell 46,674–690,June 8,2012ª2012Elsevier Inc.679structural classiﬁcation,weﬁrst queried the set of mRNA-interacting proteins against the Protein Folding Project database (Drew et al.,2011).This database provided SCOP superfamily classiﬁcations derived from sequence similarity(psi-blast),fold recognition,and Rosetta de novo structure prediction for the identiﬁed RNA-interacting proteins.An enrichment analysis of superfamilies showed an overrepresentation of folds associated with single and double-stranded RNA(dsRNA)-binding function (RNA-binding domain,eukaryotic type KH domain,and dsRNA-binding domain-like),helicases(P-loop containing nucleoside triphosphate hydrolases)and nucleases(Pin domain-like)with a corrected p value%0.05(Table S3,part A).Interestingly,we also found two structural superfamilies signiﬁcantly enriched that are associated with DNA binding(‘‘winged helix’’DNA-binding domain,and AlbA-like)suggesting that these DNA-binding folds could also interact with RNA.The‘‘winged helix’’DNA-binding domain is present in a number of RNA helicases. The AlbA-like fold was found in POP7and in C9orf23.Notably, the AlbA-like superfamily had been previously suggested to be involved in RNA binding(Aravind et al.,2003).To obtain an additional perspective of the mRNA-bound proteome associated structures,we performed Pfam and Inter-Pro domain enrichment analysis using the identiﬁed proteins. As expected,most of the signiﬁcantly enriched domains were various RNA-interaction motifs(Table S3,part B).Besides the commonly recognized RNA-binding domains,we found an overrepresentation of several domains with putative RNA-binding activity(Table S3,part B).Among these were the SWAP/SURP domain and the RAP domain,for which an RNA binding activity was suggested based on sequence compari-sons(Denhez and Lafyatis,1994;Lee and Hong,2004).The RAP domain was also signiﬁcantly enriched among the112 identiﬁed RNA interactors that were not predicted to be RNA binding(Table S3).Finally,to estimate the depth of the mRNA-bound proteome we covered in our oligo(dT)precipitations,we compared the number of identiﬁed proteins encoding at least one speciﬁc RNA-binding domain to the number of respective RNA-binding domain containing proteins observed in the deep HEK293pro-teome(Geiger et al.,2012).Figure3D shows that the majority of RNA-binding domain-containing proteins expressed in HEK293were identiﬁed by our analyses.mRNA-Bound Proteome Connects Posttranscriptional Regulation to DNA-Related ProcessesIn order to systematically examine the connectivity of the identi-ﬁed mRNA binders and their potential relationship to non-mRNA related biological processes,we generated a network based on protein-protein interaction(PPI)(Figure S3A).When comparing the PPI network of mRNA binders to a random network of equal size,we observed a higher average clustering coefﬁcient,indi-cating the presence of highly interconnected protein clusters within the network(Figure S3A).Because these clusters are indicative of functional modules mediating the regulation of complex biological processes,we analyzed the set of mRNA binders and theirﬁrst neighbors,based on PPI,for an enrichment of GO terms linked to biological processes(Ashburner et al., 2000).As expected,the most signiﬁcantly overrepresented GO terms were mRNA splicing,localization,processing,and translation(Figure S3B).In addition we observed an overrepre-sentation for DNA-related processes,namely‘‘response to DNA damage,’’‘‘DNA-dependent transcription,’’and‘‘DNA duplex unwinding’’(Figure S3B).Figure3E shows the PPI sub-network for members linked to the term‘‘response to DNA damage.’’Central to this network are XRCC6/Ku70,XRCC5/Ku80,and the DNA-activated protein kinase(PRKDC).These proteins were identiﬁed in all three pro-teomic analyses(Table S1).Besides their role in DNA double-strand break repair and recombination,the proteins have been shown to interact with RNA structures,such as the RNA-stem loop region in yeast telomerase TLC1and the RNA-component of human telomerase(Ting et al.,2005).In addition,XRCC6 had been suggested to bind internal ribosomal entry site(IRES) elements and is likely involved in the regulation of IRES-mediated mRNA translation(Silvera et al.,2006).Validation of RNA-Binding Function of SeveralCandidate mRNA InteractorsTo validate the RNA-binding activity of a subset of the identiﬁed proteins,we applied PAR-CLIP.In brief,HEK293cell lines,stably expressing epitope-tagged mRNA binders,were grown in the presence of4SU and UV irradiated.Immunopuriﬁed(IPed)and RNase-treated protein-RNA complexes were radiolabeled with T4polynucleotide kinase,separated by SDS-PAGE,and blotted onto a membrane.The radiolabeled protein-RNA complexes were visualized by phosphorimaging,while protein precipitation was monitored by western analysis.As positive controls,we usedﬁve known RNA-binding proteins:CAPRIN1(Shiina et al., 2005),HNRNPD(Knapinska et al.,2011),HNRNPR(Hassfeld et al.,1998),HNRPNU(Kiledjian and Dreyfuss,1992),as well as MYEF2,which is a transcriptional repressor(Haas et al., 1995)with an RNA recognition motif(RRM)domain.As ex-pected,the tagged proteins,IPed from UV-irradiated cells,efﬁ-ciently crosslinked to RNA,when compared to proteins that were IPed from nonirradiated cells(Figure4A).In contrast,we were unable to detect radiolabeled protein-RNA complexes in immunoprecipitations of phosphoglycerate kinase1(PGK1) and lactate dehydrogenase A(LDHA)(Figure4B),two metabolic enzymes that were not identiﬁed in our proteomic analyses as potential RNA binders(Table S1).We generated HEK293cell lines stably expressing29putative mRNA-interacting proteins as epitope-tagged versions.Twenty-one proteins could be IPed and were used in the PAR-CLIP assay(Table S4).We tested the RNA-binding activity of18candi-dates belonging to class I and three members of class II,BTF3, C16orf80,and PRDX1(Figure4C).For all proteins,except BZW1 and C16orf80,we observed an increased radioactive signal in immunoprecipitations of irradiated cells,indicating that these proteins were crosslinked to RNA,and thus are likely in close contact or directly binding to RNA.Interestingly,several of the crosslinked proteins possess enzymatic activities:ALKBH5(2-oxoglutarate oxygenase), C22orf28(RNA ligase),CSNK1E(kinase),MKRN2(ubiquitin ligase),PRDX1(peroxidase),and USP10(ubiquitin thioesterase). Furthermore several of the identiﬁed RNA-binding proteins have been implicated in transcriptional regulation either by inhibitionMolecular CellThe Protein-mRNA Interactome680Molecular Cell46,674–690,June8,2012ª2012Elsevier Inc.。

HST Imaging of the Globular Clusters in the Fornax Cluster NGC 1379

a r X i v :a s t r o -p h /9710041v 1 3 O c t 1997HST Imaging of the Globular Clusters in the Fornax Cluster:NGC 1379Rebecca A.W.ElsonInstitute of Astronomy,Madingley Road,Cambridge CB30HA,UKElectronic mail:elson@ Carl J.Grillmair Jet Propulsion Laboratory,4800Oak Grove Drive,Pasadena,CA 91109USA Electronic mail:carl@ Duncan A.Forbes School of Physics and Astronomy,University of Birmingham,Edgbaston,Birmingham B152TT,UK Electronic mail:forbes@ Mike Rabban Lick Observatory,University of California,Santa Cruz,CA 95064USA Electronic mail:mrabban@ Gerard.M.Williger Goddard Space Flight Center,Greenbelt,MD 20771USA Electronic mail:williger@Jean P.BrodieLick Observatory,University of California,Santa Cruz,CA 95064USAElectronic mail:brodie@ReceivedABSTRACTWe present B and I photometry for∼300globular cluster candidates in NGC1379,anE0galaxy in the Fornax Cluster.Our data are from both Hubble Space Telescope(HST)and ground-based observations.The HST photometry(B only)is essentially complete and free of foreground/background contamination to∼2mag fainter than the peak of the globular cluster luminosity function.Fitting a Gaussian to the luminosity function weﬁnd B =24.95±0.30 andσB=1.55±0.21.We estimate the total number of globular clusters to be436±30.To a radius of70arcsec we derive a moderate speciﬁc frequency,S N=3.5±0.4.At radii r∼3−6 kpc the surface density proﬁle of the globular cluster system is indistinguishable from that of the underlying galaxy light.At r∼<2.5kpc the proﬁle of the globular cluster systemﬂattens, and at r∼<1kpc,the number density appears to decrease.The(B−I)colour distribution of the globular clusters(from ground-based data)is similar to that for Milky Way globulars,once corrected for background contamination.It shows no evidence for bimodality or for the presence of a population with[Fe/H]∼>−0.5.Unlike in the case of larger,centrally located cluster ellipticals,neither mergers nor a multiphase collapse are required to explain the formation of the NGC1379globular cluster system.We stress the importance of corrections for background contamination in ground-based samples of this kind:the area covered by a globular cluster system(with radius∼30kpc)at the distance of the Virgo or Fornax cluster contains∼>200background galaxies unresolved from the ground,with magnitudes comparable to brighter globular clusters at that distance.The colour distribution of these galaxies is strongly peaked slightly bluer than the peak of a typical globular cluster distribution.Such contamination can thus create the impression of skewed colour distributions,or even of bimodality,where none exists.Key words:galaxies:individual:NGC1379-globular clusters:general-galaxies:star clusters1.IntroductionAn accumulating body of observations suggests that the distribution of colours of globular clusters,and by inference of their metallicities,varies signiﬁcantly from one galaxy to the next.Of particular interest is the colour bimodality which has now been observed in the globular cluster systems of several large elliptical galaxies,and suggests the presence of distinct metal rich and metal poor populations.The best example is the Virgo cD galaxy M87(NGC4486)(cf.Elson&Santiago1996).It has one population of globular clusters with colours similar to those of the Milky Way globulars,and one which is signiﬁcantly redder,with inferred metallicities∼>solar.Other galaxies whose globular cluster systems show clear bimodality include M49(NGC4472),an E2galaxy in the Virgo Cluster with the same luminosity as M87(Geisler,Lee&Kim 1996),and NGC5846,a slightly less luminous E0galaxy at the centre of a small compact group(Forbes, Brodie&Huchra1997a).Two ideas have been invoked to explain the colour bimodalities.One is that elliptical galaxies are formed during mergers in which populations of metal rich(red)globulars are created and added to a‘native’population of metal poor(blue)clusters(cf.Ashman&Zepf1992).The other is that globular cluster populations with diﬀerent mean metallicities form during a multiphase collapse of a single system(Forbes, Brodie&Grillmair1997b):metal poor globular clusters are formed early in the collapse,while metal rich ones form later,roughly contemporaneously with the stars.To understand fully the implications of the observed colour distributions for the origin of globular cluster systems,a much larger body of accurate data for systems surrounding galaxies of a variety of types and in a variety of environments is required.A question of particular importance,for example,is whether all elliptical galaxies have bimodal globular cluster systems,or whether such systems are restricted to large galaxies in rich environments.The data best suited to address this question are those acquired with the Hubble Space Telescope(HST).The resolution of HST allows even the crowded central regions of galaxies at the distance of Fornax and Virgo to be probed and allows most background galaxies to be eliminated. At these distances samples of globular clusters are complete and uncontaminated to well past the peak of the luminosity function.This paper and the others in this series are contributions to this growing database. Forbes et al.(1997a)and Grillmair et al.(1997)discuss the globular cluster systems of the Fornax cD galaxy NGC1399,its neighbour the E1galaxy NGC1404,and the peculiar galaxy NGC1316which may have undergone a recent merger.Here we present observations of the globular cluster system of NGC1379, a normal E0galaxy in the Fornax cluster.Hanes&Harris(1986)used photographic data to study the NGC1379globular cluster system toB=23.6(about a magnitude brighter than the peak of the luminosity function).They measured the proﬁle of the outer part of the system(5<r<35kpc),and estimated the total population to number∼800. More recently the globular cluster systems ofﬁve galaxies in the Fornax cluster,including NGC1379,have been studied by Kohle et al.(1996)and Kissler-Patig et al.(1997a)using V and I−band photometry obtained with the100-inch telescope at Las Campanas.Their data are50%complete at B∼24,and cover a radial range∼3−10kpc.Our observations,obtained both from the ground at the Cerro Tololo Interamerican Observatory (CTIO),and from space using HST,are described in Section2.Section3presents the results,including a caveat concerning the need for accurate background corrections for ground-based data.Ourﬁndings are summarized in Section4.2.Observations and Data ReductionIn this section we describe the HST and CTIO images upon which our results are based,and the process for detecting,selecting,and determining magnitudes for the globular cluster candidates in each case.We also discuss completeness,and contamination from foreground stars and background galaxies.At an adopted distance of18.4Mpc(m−M=31.32;Madore et al.1996),1arcsec corresponds to89pc. Reddening in the direction of the Fornax cluster is assumed to be negligible(Bender,Burstein,&Faber 1992).2.1.HST ImagingFive images of NGC1379were obtained with HST on1996March11,using the F450W(∼B)ﬁlter. (Due to technical diﬃculties,the complementary I-band images were not acquired,and are anticipated in 1997.)Three images were taken at one pointing,and two more were oﬀset by0.5arcsec.The total exposure time was5000seconds.The centre of NGC1379was positioned at the centre of the Planetary Camera (PC)chip to aﬀord the greatest resolution in the most crowded regions.Details of the reductions are given by Grillmair et al.(1997).Brieﬂy,the images were reduced using the standard pipeline procedure.TheVISTA routine SNUC was used toﬁt and subtract the underlying galaxy.We ran DAOPHOT II/ALLSTAR (Stetson1987)separately on the sum of theﬁrst three images and the last two images,requiring that detections appear in both lists to qualify as real objects.We adopted a detection threshold of3σand measured magnitudes byﬁtting a point-spread function(PSF).Extended objects were eliminated by visual inspection.Count rates were converted to B magnitudes using the gain ratios and zeropoints given by Holtzman et al.(1995;1997,private communication).Photometry is available on request from CJG.Figure1shows a mosaic of the four WFPC2chips,with the galaxy subtracted.The total area of the ﬁeld,excluding two60pixel wide unexposed borders on each chip,is4.8arcmin2.The scale of the Wide Field Camera(WFC)is0.0996arcsec pixel−1,and of the PC,0.0455arcsec pixel−1.At the distance of NGC1379,one WFC pixel corresponds to∼9pc and one PC pixel to∼4pc.A globular cluster with size typical of those in the Milky Way(core radius∼2pc,half-mass radius∼10pc,and tidal radius∼50pc) will thus appear essentially unresolved in our images.A total of∼300objects were detected and measured in ourﬁeld.To determine the completeness of this sample,3000artiﬁcial PSFs(100at each of30magnitude levels)were added to the images,and the images were then processed in a manner identical to that for the original data.The completeness of the sample as a function of magnitude is shown in Figure2.The sample is∼100%complete to B=26for the WFC chips(80%of the sample),and to B=25.5for the PC chip.At B>26the completeness begins to drop rapidly.Photometric errors areδB≈0.10mag at B∼<25,rising to0.15mag at B=26.Next we consider the extent to which our sample may be contaminated by foreground stars and background galaxies.Few foreground stars are expected in an area of only4.8arcmin2at this Galactic latitude,and most background galaxies are resolved and thus easily distinguished from globular clusters. The main source of contamination is compact,spherical background galaxies.To determine the expected level of contamination in our sample,we observed a backgroundﬁeld located∼1.4degrees south of the center of the Fornax cluster.The exposure time was5200seconds,so the limit of detection is comparable to that for the NGC1379sample.The image was processed and the sample selected in the same way as for the NGC1379ﬁeld(see Grillmair et al.1997).Figure3shows a colour-magnitude diagram(CMD)for the84unresolved objects detected in the backgroundﬁeld.The sample becomes incomplete at B>26.5,but as we shall see,this is∼1.5magnitudes fainter than the peak of the luminosity function,and so will not aﬀect our results.At B∼<26.5the objects have a wide range of colours,with the majority concentrated around(B−I)∼1.The B luminosityfunction for the background sample is plotted in Fig.4,which also shows the luminosity function for the ∼300candidate globular clusters.The background luminosity function is tabulated in Table1.Since these background number counts are applicable to HST studies of any unresolved population at high latitude, we also include in Table1the I-band luminosity function for the backgroundﬁeld.This is plotted as the solid histogram in Fig. 5.As a check on the consistency of the sample,and to assess the amplitude of spatialﬂuctuations in the background on this scale,we compared the luminosity function in Fig.5with the Medium Deep Survey(MDS)star count data from17high latitudeﬁelds obtained with HST in the V and I bands(Santiago,Gilmore&Elson1996).The dotted histogram in Fig.5shows the I-band luminosity function for the MDS data,normalized to an area of4.8arcmin2.The two distributions are in excellent agreement to I≈23.5which is the limiting magitude of the MDS star count data.Fainter than this stars can no longer be reliably distinguished from compact galaxies.The MDS stellar luminosity function in V, normalized to the same area,is also included in Table1.Finally,Fig.5also shows a luminosity function for stars and galaxies from the Canada-France Redshift Survey(CFRS)(Lilly et al.1995;S.Lilly1997,private communication).The magnitudes were measured in an aperture of diameter3arcsec.The I-band data cover an eﬀective area of∼425arcmin2,and have again been normalized to an area of4.8arcmin2.These data illustrate the much larger degree of background contamination to be expected in the extreme case where no galaxies(with I∼>21)can be distinguished morphologically from unresolved objects.Background contamination in typical ground-based samples of globular clusters in distant galaxies will fall somewhere between the CFRS and the HST histograms.2.2.Ground–based ImagingBroadband B and I images of NGC1379were taken with the CTIO1.5m telescope,using a Tek2048 x2048array with a pixel scale of0.44arcsec pixel−1.Total exposure times were9000and3900seconds for the B and I images respectively.Although the seeing was only∼1.5arcsec,conditions remained photometric throughout the night of1995December24.Reduction was carried out in the standard way(i.e. bias and dark subtraction,ﬂat–ﬁelding and sky subtraction).After combining,the images were calibrated using aperture photometry from the catalogs of of Longo&de Vaucouleurs(1983)and de Vaucouleurs& Longo(1988).This procedure gave an rms precision of better than0.05mag.For the selection of globular cluster candidates in the CTIO images,we employed an iterative procedure using DAOPHOT II.We measured the background noise in both images and set the threshold for singlepixel detection at5σ,i.e.ﬁve times the noise due to the background.The other important detection parameters,SHARPness and ROUNDness(designed to weed out extended objects and cosmic rays),were initially given a large range.For each detected object we measured a3pixel radius aperture magnitude and applied an aperture correction based on a curve-of-growth analysis for a dozen isolated globular clusters. The rms error in the aperture correction is∼0.05mag.We then compared our B-band candidate list with the positions and B magnitudes of globular clusters detected with HST’s WFC.Our candidate list was matched to the HST list with the condition that the HST globular cluster lie within3CTIO pixels of our object.With this condition,28objects were matched. The average magnitude diﬀerence between the CTIO and HST B magnitude is0.03mag.Such excellent agreement is gratifying from a photometric standpoint,and reassures us that we have matched the data sets correctly.The matched clusters have SHARPness parameters0.3to0.7and ROUNDness−0.4to0.4. Assuming that these globular clusters are representative of all globular clusters in the CTIO image,we re–ran DAOPHOT II with the new restricted range in SHARPness and ROUNDness parameters.This resulted in the exclusion of about half the objects in the original sample,which are presumably background galaxies.The same parameters were used for both the B and I images.Figure6shows a CMD for365 objects selected in this way.A comparison of the CTIO object list and the HST list,within the area in common,will also indicate how complete our detection is as a function of magnitude.The completeness function estimated this way is shown in Fig.7.Our sample is∼100%complete at B∼21and∼50%complete at B∼23.5.The photometric errors are<0.1mag for all magnitudes brighter than our50%completeness limit.At this point it is necessary to investigate the level of contamination of our ground-based sample by foreground stars and background galaxies.Since theﬁeld is much larger than that observed with HST(211 arcmin2compared to4.8arcmin2),and since the resolution is not suﬃcient to distinguish many background galaxies from unresolved sources,contamination from both stars and galaxies will be signiﬁcant.Since no background comparisonﬁeld was obtained,we rely instead on the CFRS data in the B and I-bands, discussed above,to estimate a correction for contamination.Figure8shows histograms of(B−I)for the full NGC1379sample in Fig.6,and for the CFRS sample of objects with19<B<23.Since selection eﬀects in the two samples are diﬀerent,we do not normalize the CFRS sample directly using the known area of the survey.Rather,we normalize it to match the tail of objects with(B−I)>2.5in Fig.8,whose colours are too red to be globular clusters.From this wecan infer the relative number of contaminants with(B−I)<2.5.The CFRS sample is strongly peaked at (B−I)∼1.0,which is just∼0.4mag bluer than the peak of the NGC1379sample.The NGC1379sample is lacking many of the bluest objects in the CFRS sample;this is probably because fainter galaxies are on average bluer,and the CFRS sample is much more complete than the NGC1379sample at the faintest magnitudes.Also,our process of excluding galaxies from the CTIO sample on the basis of ROUNDness would preferentially eliminate edge-on spirals,which are systematically bluer than ellipticals.The most important point illustrated by Fig.8is that in the ground-based data,much more so than in the HST data,failure to properly correct for foreground/background contamination may lead to a signiﬁcant error in the deduced colour of the peak of the globular cluster colour distribution,and possibly to the erroneous impression of a bimodal colour distribution.With ground-based data obtained with better seeing it would of course be possible to exclude a greater proportion of background galaxies,so that any skewing of the colour distribution would be less severe.3.Properties of the NGC1379globular cluster systemThe principal properties of a globular cluster system which any theory of its origin and evolution must account for are the luminosity function(in particular,the absolute magnitude of its peak,and the width of the distribution),the colour distribution(in particular,the colour of the peak and the presence or absence of any bimodality),the radial distribution,and the total number of clusters,N tot,and therefore the speciﬁc frequency,deﬁned as S N=N tot100.4(M V+15),where M V is the absolute magnitude of the galaxy.With B tot=12.07(Tully1988)and(B−V)=0.89(Faber et al.1989)we infer an apparent magnitude for NGC 1379of V tot=11.18.With distance modulus31.32,this implies M V=−20.16.This is the integrated magnitude within a radius of∼70arcsec.3.1.Luminosity functionThe luminosity function for the NGC1379globular cluster candidates derived from the HST datais shown in Fig. 4.This sample includes∼300unresolved objects and is essentially complete and uncontaminated well past the peak.The background-corrected luminosity function,that is,the diﬀerencebetween the solid and dashed histograms in Fig.4,is shown in Fig.9.Fitting a Gaussian function to the luminosity function in Fig.4,using a maximum-likelihood technique which is independent of binning, and takes into account completeness,background contamination,and photometric errors(Secker&Harris 1993),we derive a peak magnitude of B =24.95±0.30,and a widthσB=1.55±0.21.This Gaussian is shown as the solid curve in Fig.9.We know of no other direct measurements of B for this system to compare with ours.However,our values are the same,to within the errors,as the values measured for the globular cluster systems of two other Fornax galaxies,NGC1399and NGC1404(Grillmair et al.1997).Also,for four elliptical galaxies in the Virgo cluster Harris et al.(1991)ﬁnd an average value B =24.77±bining this with the value(m−M)F ornax−(m−M)V irgo=0.08±0.09(Kohle et al.1996)implies B =24.85±0.22for NGC 1379.This is in good agreement with our value.We may either use the measured peak magnitude to infer a distance modulus for NGC1379,assuming a‘universal’value for the absolute magnitude of the peak,or we may adopt a distance modulus and infer an absolute magnitude.While the value of M V has been measured for many globular cluster systems, there are fewer B−band studies.Sandage&Tamman(1995)quote M B =−6.93±0.08for the Milky Way and M31globular clusters,while Ashman,Conti&Zepf(1995)give M B =−6.50for the Milky Way clusters.Adopting the Cepheid distance for NGC1379and our measured peak magnitude,implies M B =−6.37±0.36,which is in good agreement with the Ashman et al.value and somewhat fainter than the Sandage&Tamman value.As we shall see below,the(B−I)distribution of NGC1379globular clusters appears very similar to that for the Milky Way globular clusters.It should therefore be safe to assume that the(B−V)distribution is also similar.Adopting B−V =0.7for the Milky Way globulars(Harris1996),we may convert our value for the peak magnitude of the luminosity function, B =24.95±0.30,to V =24.25±0.30.This value is somewhat fainter than the value found by Kohle et al.(1996)of V =23.68±0.28,but their data reach just to the peak of the luminosity function,and the errors in their faintest bin are large.The sense of the diﬀerence is consistent with our result above thatﬁtting only the brighter part of the luminosity function results in a peak magnitude which is too bright.With the Cepheid distance modulus of31.32,our V value implies M V =−7.07±0.36,which is typical for elliptical galaxies of this absolute magnitude (Harris1991).Our value ofσB=1.55±0.21for the NGC1379globular cluster system is larger than the valuesσB=1.07and0.89quoted by Sandage&Tamman for the Milky Way and M31respectively,but is consistent with the valuesσB=1.37±0.07andσB=1.39±0.12found by Grillmair et al.(1997)for the globular cluster systems of NGC1399and NGC1404respectively.It is also consistent with the value σB=1.46±0.07for four elliptical galaxies in the Virgo cluster(Harris et al.1991).3.2.Radial proﬁleHarris&Hanes(1987)compared the radial proﬁle of the NGC1379globular cluster system with the surface brightness proﬁle of the galaxy itself from Schombert(1986),over the radial range5−35kpc. They detected no diﬀerence between the two,although their uncertainties were large.Kissler-Patig et al. (1997a)found the same result from their ground-based data over the range3−10kpc.Radial surface density proﬁles for our HST sample of NGC1379globulars are shown in Figs.10a and b.Corrections have been made to compensate for the fraction of each annulus which falls outside theﬁeld of view of the WFPC2.In Figure10a proﬁles are plotted for both the complete,uncontaminated sample with B<25.5, and for the objects with B>26.5(with no correction for completeness),which are expected to be primarily background galaxies.Indeed,the fainter sample shows almost no radial gradient.The dashed line is the background level(∼9.0objects per arcmin2)measured from the backgroundﬁeld for B>26.5.The agreement is excellent.The background level for the brighter sample,again measured from the background ﬁeld,is2.7objects per arcmin2,and is shown as the solid line.The diﬀerence between the radial proﬁle and the background level is probably due to small numbers,but may indicate a small amount of residual contamination:theﬁve outermost points represent an average of<4objects each,and the total number of objects in the backgroundﬁeld with B<25.5is just13.The radial proﬁle of the brighter sample decreases smoothly from∼10−80arcsec(∼1−7kpc),at which point the globular cluster system is lost in the background.Figure10b shows the radial proﬁle for the sample with B<25.5,with a background of2.7objects per arcmin2subtracted.Superposed is a curve representing the surface brightness proﬁle of the underlying galaxy from Kissler-Patig et al.(1997a),scaled arbitrarily to match the proﬁle of the globular cluster system.The two proﬁles agree well at∼35−70 arcsec(3–6kpc).There is no evidence that the surface brightness proﬁle of the globular cluster system is shallower than that of the underlying galaxy light out at least to the limit of our data at r∼7kpc.The logarithmic slope of the proﬁle in Fig.10b is−2.4at r∼>35arcsec.Inwards of r∼30arcsec(∼2.5kpc)the proﬁle of the globular cluster systemﬂattens out.This core structure seems to be a common feature of the globular cluster systems of elliptical galaxies,and the radius at which theﬂattening occurs correlates with the galaxy luminosity(Forbes et al.1996).The mean surface density within∼10arcsec of the centre of NGC1379is∼200±60clusters per arcmin2.The core radius, where the surface density has fallen to half its central value,is r c∼23±6arcsec(2.0±0.5kpc).This is consistent with values for other galaxies with absolute magnitude comparable to that of NGC1379.Such a core structure is not present in the underlying galaxy light,which,while it changes slope slightly at r∼50 arcsec,rises with constant slope inwards to at least10arcsec(Schombert1986).Inwards of∼10arcsec(∼1kpc),the surface density of globular clusters appears to decrease.This radius corresponds to220pixels on the PC,and as can be seen from Fig.1,crowding even in these inner regions is not severe.Extrapolating a smoothly rising proﬁle to the center of the galaxy would require6 clusters instead of2in the innermost bin,and14instead of12in the second bin.If the radial distribution really were steadily rising,this would imply that our data are only∼70%complete in these combined bins. Closer analysis of our completeness tests reveals that we are in fact97%complete for B<25.5in the region r<220pixels(the slight reduction over the PC-averaged completeness value being a consequence of the increased noise due to the integrated stellar light of the central regions of the galaxy).We conclude that the central dip in the cluster surface density distribution is not an artifact of our analysis,and may indeed indicate a quite substantial drop in the volume density of clusters near the nucleus of the galaxy.Radii less than1kpc from the nucleus are where we expect tidal stresses to begin to take a toll on the numbers of globular clusters(Lauer&Kormendy1986,Grillmair,Pritchet,&van den Bergh1986),either during their formation or subsequently through tidal stripping(Grillmair et al.1995),particularly if the clusters are on box orbits.We now turn to the ground-based data which cover a much greater area than the HST data.The radial distribution for the full sample from Fig.6is shown in Fig.11.The surface density for the CTIO sample has been increased by0.74in log N to match the HST sample,which is shown as the solid curve.At r∼>100arcsec,the proﬁle of the CTIO sample is essentiallyﬂat,suggesting that the sample is composed overwhelmingly of foreground/background objects at these radii.Even a colour selection is unlikely to help disentangle the background contamination since,as shown in Fig.8,background galaxies have a similar colour distribution to the full NGC1379sample.We therefore conclude that,despite the larger spatial coverage of the ground-based sample,in the absence of a suitable background calibrationﬁeld it does not contribute much to our knowledge of the radial structure of the NGC1379globular cluster system notcovered by our HST data.It is also unable to probe the innermost regions of the cluster system,due to crowding.3.3.Colour distributionIn the absence of an I-band image from HST,we attempt to extract a sample of globular clusters from the ground-based data which is as uncontaminated as possible,to investigate the(B−I)colour distribution.One approach is to select only objects with r<100arcsec,since these show a radial gradient in surface density(Fig.11).This gives a sub-sample of35objects.The background level inferred from the radial distribution of the CTIO sample at r>100arcsec is1.6arcmin−2,implying that14of the35 globular cluster candidates at r<100arcsec,or40%of this sample,are background galaxies or foreground stars.The colour distribution for the35objects at r<100arcsec,along with those for the full CTIO sample and for the CFRS sample,is shown as the dashed histogram in Fig.8.The r<100arcsec sample clearly peaks at a redder colour than the full sample,underlining the fact that the colour distribution of the uncorrected sample is misleading.Figure12shows the same histogram for the r<100arcsec sample,along with the histogram of residuals obtained by subtracting the normalized CFRS histogram from the full CTIO histogram in Fig.8.A histogram for95globular clusters in the Milky Way(Harris1996)is also shown.The colour distributions of the NGC1379sample with r<100arcsec,and with the CFRS sample subtracted,have a very similar peak colour,(B−I)≈1.6,which is indistinguishable from that for the globular cluster system of the Milky Way.Using the relation between(B−I)colour and metallicity from(Couture et al.1990)we infer from the peak colour a metallicity of[Fe/H]∼−1.5.Forbes et al.(1997b)plot a relation between mean metallicity of globular clusters and parent galaxy magnitude for11galaxies with−21<M V<23,with bimodal globular cluster colour distributions.Theyﬁnd that,while in the metal rich populations([Fe/H]>−0.5)there is a strong correlation between mean globular cluster metallicity and galaxy magnitude,for the metal poor populations the scatter is much greater and the correlation much weaker.There is,however,a trend for less luminous galaxies to have globular clusters with a lower mean metallicity,and our results for NGC1379are consistent with this trend.There is a tail of bluer objects in the NGC1379samples which is probably comprised of residual。

The Halpha emission of the spiral galaxy NGC 7479

a r X i v :a s t r o -p h /9906390v 1 24 J u n 1999The H αemission of the spiral galaxy NGC 7479Almudena Zurita (azurita@ll.iac.es ),Maite Rozas(mrozas@ll.iac.es )and John E.Beckman (jeb@ll.iac.es )Instituto de Astrof´ısica de Canarias,38200-La Laguna,Tenerife,SPAINAbstract.We use the catalogue of H ii regions obtained from a high quality continuum-subtracted H αimage of the grand design spiral galaxy NGC 7479,to construct the luminosity function (LF)for the H ii regions (over 1000)of the whole galaxy.Although its slope is within the published range for spirals of the same morphological type,the unusually strong star formation along the intense bar of NGC 7479prompted us to analyze separately the H ii regions in the bar and in the disc.We have calculated the physical properties of a group of H ii regions in the bar and in the disc selected for their regular shapes and absence of blending.We have obtained galaxy-wide relations for the H ii region set:diameter distribution function and also the global H αsurface density distribution.As found previously for late-type spirals,the disc LF shows clear double-linear behaviour with a break at log L Hα=38.6(in erg s −1).The bar LF is less regular.This reﬂects a physical diﬀerence between the bar and the disc in the properties of their populations of regions.1.Observations,data reduction and production of the H ii region catalogue.The observations were made through the TAURUS camera on the 4.2m William Herschel Telescope on La Palma.The detector used was an EEV CCD 7with projected pixel size 0”.279×0”.279.Observing conditions were good,with photometric sky and 0”.8seeing.After standard reduction routines,we obtained the calibrated H αcontinuum-subtracted image (Fig.1)by subtracting an image through a non-redshifted ﬁlter to the image obtained through a 15˚A ﬁlter with central wavelength equal to the redshifted H αemission from the galaxy and calibrating with observations of standard stars.As a selection criterionTable I.NGC 7479:basic parameters (data from RC3catalogue,except P.A.and i ,both from Laine &Gottesman 1998)to construct the H ii region catalogue we speciﬁed that a feature must2Figure1.Representaction of the Hαcontinuum-subtracted image of NGC7479 contain at least nine contiguous pixels,each with an intensity of at least three times the r.m.s noise level of the local background.The r.m.s.noise of the background-subtracted Hαimage is15instrumental counts,which means that lower limits to the luminosity of the detected H ii regions,and to the radius of the smallest catalogued regions are, respectively,log L Hα=37.65erg s−1and≈75pc.The detection and cataloguing of the H ii regions were performed using a new pro-gram developed by C.Heller.The program identiﬁes each H ii region, measures the position of its centre,derives the area in pixels and the ﬂux of each region,integrating all the pixels belonging to the region and subtracting the local background value.We catalogued1009H ii regions in NGC7479,and for all the H ii regions we determined equa-torial coordinate oﬀsets from the nucleus and deprojected distances to the centre(in arcsec)using the inclination and position angles given by Laine&Gottesman(1998)(i=51◦,PA=22◦).In Fig.2we show schematically the positions of the H ii regions in the disc of NGC7479, on a deprojected RA-dec grid centred on the nucleus of the galaxy.3Figure2.Representation of the positions of the measured H ii regions.Symbols show ranges of log L.Coordinates of the centre of the image are R.A.=23h4m 56.64s Dec=12◦19′22.9′′(J2000)2.Luminosity functionsThe full count of H ii regions seen in Fig.2is presented in diﬀerential form,as a log-log plot in bins of0.15dex,in Fig.3.The luminosity range in Hαis limited at the high end by the luminosity of the brightest detected region,and at the low end by our criteria for a detection.The limit of completeness is log L Hα=38.0,so that the apparent broad peak in the distribution just below this luminosity is an artefact due to observational selection.The best single linearﬁt of the data to the points above log L Hα=38,corresponding to a power lawdN(L)=ALαdL,has a gradient-0.83±0.06(i.e.α=-1.83±0.06).We have also carried out a bi-linearﬁt,based on the evidence that at log L∼38.8there is an apparent change in slope in the LF and we explained it as a manifestation of a transition in the physical properties of the regions, but in those cases the luminosity of the transition was always log L Str∼38.6(∼0.2dex lower than in NGC7479).As NGC7479is a barred4galaxy with unusually strong star formation in the bar,we tested the hypothesis that the diﬀerences in the LF might be due to this intense star formation activity under somewhat diﬀerent physical conditions from those in the disc.To perform this test we constructed separately the LF’s for the H ii regions of the bar and of the disc.The results are shown in Fig.3and the LF’s slopes are in Table II.In the LF for the629H ii regions detected in the disc(Fig.5)the change in slope is cleaner than for the total LF and occurs at log L=38.65±0.15(in erg s−1).The change of slope,accompanied by a slight bump in the LF,has been detected in the seven galaxies so far examined in this degree of detail(see e.g.Rozas,Beckman&Knapen1996,Knapen et al.1993).Adding the value for NGC7479to the set of values previouslyFigure3.Luminosity function for the complete sample of H ii regions from the cata-logue with the best linear(a)and bi-linearﬁt(b),based on the previous experience with late-type spirals.(c)shows the bi-linearﬁt for the luminosity function of the H ii regions of the disc;the change in slope is seen at log L Hα=38.65.The luminosity function of the bar(d)is clearly less well-behaved than that for the disc.measured for other galaxies,weﬁnd that the rms scatter in L Str in the5 full set of objects is0.08mag.This low scatter can be explained if the IMF at the high luminosity end of the mass function changes little from galaxy to galaxy,i.e.varies little with metallicity.Another necessary condition is that the rate of emission of ionizing photons from a young stellar cluster rises more rapidly than the mass of its placental cloud, a condition which we have examined observationally(Beckman et al. 1999),and shown to hold.If we look at Fig.3,where the LF for the 380H ii regions of the bar is presented,we can see clearly that the irregularity found in the total LF is due to the incorporation of the H ii regions of the bar in the total LF.Clearly star formation conditions in the bar diﬀer from those in the disc.Table II.Luminosity limits,LF slopes and correlation coeﬃcients for thetwo luminosity ranges(in erg s−1)for the LF for the complete sample ofH ii regions from the catalogue and for the LF of the H ii regions of thedisc.Total LF single linearﬁt38.1≥log L≤40.5-1.83±0.060.969bi-linearﬁt38.1≥log L≤38.8-1.33±0.040.97938.8≥log L≤40.5-1.99±0.070.9783.Other Statistical properties.−The integral diameter distribution(number of regions with diameters greater than a given value as a function of diameter) is given in Fig.4a From our own and other published studies it has been found that the diameter distribution of the H ii re-gions can be wellﬁtted by an exponential of form N(>D)= N o exp(−D/D o)where D o is a characteristic diameter,and N o is an(extrapolated)characteristic value for the total number of regions.From the observations represented in Fig.4a,we obtain the values N o=5000±300and D o=110±2pc(with H o=75km s−1Mpc−1).D o has a value predicted for a galaxy of its mea-sured absolute luminosity according to the observational plot ofD o versus luminosity(Hodge1987).−Theﬂux density distribution is shown in Fig.4b.It was found by dividing the disc into rings using the values of P.A.and i of6diameter distribution function of all H ii regions of NGC7479.The straight line indicates the bestﬁt.b.)Flux density distribution of all the H ii regions of NGC7479as a function of the deprojected distance to the centre of the galaxy.the galaxy(Table I).There is a clear trend to lowerﬂux,with increasing radius,which clearly reﬂects the standard radial decline in surface density of each major component of this disc galaxy.This underlying radialﬂux density can beﬁtted by an exponential of form f(r)=B exp(−r/h Hα)from which we derive a value for the Hαscale length,h Hα=(2.4±0.3)kpc.4.Physical properties.Using a distance to NGC7479of31.92Mpc we employed the standard theoretical formula(Spitzer1978)relating the surface brightness of an H ii region with its emission measure(Em),(assuming case B recombination and T∼104K)to calculate the Em of the H ii regions. We performed this for a total of39regions,(22in the disc and the rest in the bar),covering the full range of observed radii and chosen as isolated to minimize the uncertainties in calculating their luminosities due to the overlapping.Results are shown in Fig.5a.In Fig.5b.we show the rms electron density(derived from the Em, <N e>rms=7agree well with those found by Kennicutt 1984and by ourselves (Rozas,Knapen &Beckman 1996)for extragalactic H ii regions.Due to ob-servational selection,these tend to be more luminous and larger than Galactic regions.We also computed the ﬁlling factors,using δ=(<N e >rms /N e )2.The implicit model is that an H ii region is internally clumpy,so that the observed ﬂux comes from a high density component,which occupies a fraction δ(ﬁlling factor)of the total volume;the rest of the volume is ﬁlled with low density gas which makes a negligible contribution to the observed emission line strengths.To calculate δ,we need to know the in situ N e for each region.We have not measured these values for NGC 7479,but have used a “canonical”mean value of 135cm −3obtained by Zaritzky et al.1994for 42H ii regions in a large sample of galaxies.The ﬁlling factors for the regions range from 4.3×10−4to 1.5×10−3,a range which coincides well with those found for 4galaxies in Rozas et al.1996b.Values of <N e >rms can also be used to estimate the mass of ionized gas,which range from 3000M ⊙to 1.5×106M ⊙.These physical properties,the emission of Lyman cotinuum photons and the equivalent number of O5V stars required to supply the luminosity of the regions are given in a more detailed paper (Rozas,Zurita &Beckman 1998).1002003004001002003004003456Figure 5.Emission measure versus radius (a.),rms electron density versus radius (b.)for the 39selected H ii regions in NGC 7479.85.ConclusionsThe main results of the present work,where we analyze the statistics and properties of H ii regions in NGC7479are summarized below.−Using a high quality continuum-subtracted Hαimage of the grand-design spiral NGC7479,we have catalogued1009H ii regions.−The slope of the LF agrees broadly with slopes for other galaxies of comparable morphological types.−We have found a change in slope in the LF of the H ii regions of NGC7479that occurs at a luminosity slightly higher(≃0.2 dex)than that found in other galaxies of the same morphological type.Due to the intense star formation in the bar of NGC7479, we decided to construct separately the LF’s for the H ii regions of the bar and of the disc,ﬁnding:1.the LF of the disc is in no way diﬀerent from that found inprevious papers for galaxies of the same morphological type,2.The anomaly in the global LF is thus due to diﬀerent star for-mation conditions in the bar.This must be due to the eﬀects ofgas dynamical parameters on the stellar IMF,and the physicalconditions in the clouds.−The integrated distribution function of the H ii region diameters can be wellﬁtted by a exponential function.−The densities,ﬁlling factors and masses derived from the luminosi-ties and sizes of a selected set of representative regions,through the range of observed luminosities for NGC7479,are in agreement with those found in the previous literature on extragalactic H ii regions.ReferencesBeckman,J.E.,Rozas,M.,Zurita,A.&Knapen,J.H.1999,AJ,,submittedde Vaucouleurs,G.,de Vaucouleurs,A.,Corwin,H.G.,Buta,R.J.,Paturel,G., Fouqu´e,P.,1991Third Reference Catalogue of Bright Galaxies(RC3),Springer, New YorkHodge,P.W.1976,ApJ,205,728Knapen,J.H.,Arnth-Jensen,N.,Cepa,J.,Beckman,J.E.1993,AJ,106,56 Laine S.&Gottesman S.T.1998,MNRAS,297,1041Osterbrock,D.E.1974,Astrophysics of gaseous nebulae,Freeman,San Francisco Rozas,M.,Beckman,J.E.&Knapen,J.H.1996a,AA,307,735Rozas,M.,Knapen,J.H.&Beckman,J.E.1996b,AA,312,275Rozas,M.,Zurita,A.,Heller,C.H.&Beckman,J.E.,1998,AAS,135,145 Spitzer,L.,1978,Physical Processes in the Interstellar Medium,Wiley,New York Zaritzky,D.,Kennicutt,R.C.&Huchra,J.P.1994,ApJ,420,87。

Dark Matter Halos around Galaxies

Dark Matter Halos around GalaxiesP.SalucciSISSA,via Beirut2-4,I-34013Trieste,ItalyM.PersicOsservatorio Astronomico di Trieste,via Beirut2-4,I-34013Trieste, ItalyAbstract.We present evidence that all galaxies,of any Hubble type and lumi-nosity,bear the kinematical signature of a mass component distributed diﬀerently from the luminous matter.We review and/or derive the DM halo properties of galaxies of diﬀerent morphologies:spirals,LSBs,ellip-ticals,dwarf irregulars and dwarf spheroidals.We show that the halo density proﬁleM h(x)=M h(1)(1+a2)x3x2+a2(with x≡R/R opt),across both the Hubble and luminosity sequences, matches all the available data that include,for ellipticals:properties of the X-ray emitting gas and the kinematics of planetary nebulae,stars, and HI disks;for spirals,LSBs and dIrr’s:stellar and HI rotation curves; and,ﬁnally,for dSph’s the motions of individual stars.The dark+luminous mass structure is obtained:(a)in spirals,LSBs, and dIrr’s by modelling the extraordinary properties of the Universal Rotation Curve(URC),to which all these types conform(i.e.the URC luminosity dependence and the smallness of its rms scatter and cosmic variance);(b)in ellipticals and dSph’s,by modelling the coadded mass proﬁles(or the M/L ratios)in terms of a luminous spheroid and the above-speciﬁed dark halo.A main feature of galactic structure is that the dark and visible matter are well mixed already in the luminous region.The transition between the inner,star-dominated regions and the outer,halo-dominated region,moves progressively inwards with decreasing luminosity,to the extent that very-low-L stellar systems(disks or spheroids)are not self-gravitating,while in high-L systems the dark matter becomes a main mass component only beyond the optical edge.A halo core radius,comparable to the optical radius,is detected at all luminosities and for all morphologies.The luminous mass fraction varies with luminosity in a fashion common to all galaxy types:it is comparable with the cosmological baryon fraction at L>L∗but it decreases by more than a factor102at L<<L∗.1LSBsFigure1.The loci populated by the various families of galaxies inthe central brightness vs.luminosity plane.For each Hubble type,the central halo density increases with de-creasing luminosity:sequences of denser stellar systems(dwarfs,ellipti-cals,HSBs,LSBs in decreasing order)correspond in turn to sequences of denser halos.Then,the dark halo structure of galaxiesﬁts into a well ordered pattern underlying a uniﬁed picture for the mass distribution of galaxies across the Hubble sequence.1.IntroductionOver scales∼1kpc up to the Hubble radius,the dynamics of cosmological systems is inﬂuenced,and often dominated,by non-radiating matter which re-veals itself only through a gravitational interaction with the luminous matter. As the observational evidence has been accumulating,it has become apparent that understanding the nature,history and structural properties of this dark component,is the focal point of Cosmology at the end of the millennium.In particular,the halos of dark matter,detected around galaxies,have driven the dissipative infall of baryons that,modulo a variety of initial conditions,has built the bulge/disk/spheroid systems we observe today.It is remarkable that the∼1011galaxies present within the Hubble radius can be classiﬁed in a very small number of types:ellipticals,spirals,low-surface-brightness(LSB)galaxies, dwarf spirals and dwarf irregulars.The main characterizing property of these families is their position in theµ0,M B(central surface brightness,magnitude)2plane(see Fig.1).In this plane spiral galaxies lie at the center and show a very small range(∼0.5mag)inµ0;ellipticals are very bright systems,and span only a factor10in luminosity,that however well correlates with central brightness; LSB galaxies are the counterpart of spirals at low surface brightness;and,ﬁnally, dwarfs are very-low-luminosity spheroids or disks which barely join the faintest normal systems and span the largest interval inµ0.Cosmologically,all galaxy types are equally important for at least two reasons:(a)those types having a lower average luminosity are however much more numerous and hence can store a signiﬁcant amount of baryons,and(b)the properties that characterize and diﬀerentiate the various Hubble types,i.e.the angular momentum content and the stellar populations,are intimately related to the process of galaxy formation.A systematic presence of dark matter wasﬁrst found in spirals,speciﬁcally from the non-keplerian shape of their rotation curves(Rubin et al.1980;Bosma 1981),and in dSph’s from their very high tidal M/L ratios(Faber&Lin1983). Later,dark matter has been sistematically found also around dwarf spirals and LSB galaxies(Romanishin et al.1982).For ellipticals,the situation is less clear:certainly at least some(if not all)show evidence of a massive dark halo in which the luminous spheroid is embedded(e.g.Fabricant&Gorenstein1983). Therefore,the claim for the ubiquitous presence of dark halos around galaxies may be observationally supported.Theoretically,this claim is the natural outcome of the bottom-up cosmo-logical scenarios in which galaxies form inside dark matter halos(probably with a universal density proﬁle;e.g.,Frenk et al.1988and Navarro et al.1996). We point out that this prediction has so far been tested only in spirals where reliable DM proﬁles have been obtained(Persic,Salucci&Stel1996;hereafter PSS96).Along the Hubble sequence,the systematic presence of dark matter in galaxies and its relation with the luminous matter has so far been poorly known. However,in the past year or so a number of observational breakthroughs(some of which presented at this conference)have allowed us to obtain the gravita-tional potential in numerous galaxies of diﬀerent Hubble types(including also LSB galaxies).The time is now ripe for attempting to derive the general mass proﬁle of dark matter halos,as a function of galaxy luminosity and morphology.In this paper(which is also a review describing much recent work),we will try to answer, for theﬁrst time,a simple(cosmological)question:Given a galaxy of Hubble type T and luminosity L,which halo is it embedded in?The aim of this article is then to derive/review,by means of proper mass modelling,the mass distribution in galaxies of diﬀerent luminosities and Hubble Types and toﬁt all of the various pieces into one uniﬁed scheme of galaxy structure.Notice that the theoretical implications of the results presented here will be discussed elsewhere.In detail,the plan of the paper is as follows:in section2we review the properties of DM halos in spiral galaxies;in section 3we work out the DM mass distribution in LSBs and perform a comparative analysis with that in spirals;in section4we derive/review the halo properties of elliptical galaxies;in sections5and6we derive the properties of dwarf galaxies, irregulars and spheroidals.Finally,in section7we propose a uniﬁed scheme for the DM halos around galaxies and their interaction with the luminous matter.(A value of H0=75km s−1Mpc−1is assumed throughout.)3Figure2.The RC slope at R opt vs luminosity(left)and V opt(right)for the1100RCs of Persic&Salucci1995and PSS96.2.Spiral GalaxiesThe luminous(∼stellar)matter in spiral galaxies is distributed in two com-ponents:a concentrated,spheroidal bulge,with projected density distribution approximately described byI(R)=I0e−7.67(R/r e)1/4(1) (with r e being the half-light radius;but see Broeils&Courteau1997),and an extended thin disk with surface luminosity distribution very well described(see Fig.3)by:I(R)=I0e−R/R D(2) (with R D being the disk scale-length;Freeman1970).Let us take R opt as the ra-dius encircling83%of the integrated light:for an exponential disk R opt=3.2R D is the limit of the stellar disk.The relative importance of the two luminous com-ponents deﬁnes the Hubble sequence of spirals,going from the bulge-dominated Sa galaxies to the progressively more disk-dominated Sb/Sc/Sd galaxies.We recall that the spiral arms are non-axisymmetric density perturbations,traced by newly-formed bright stars or HII regions,which are conspicuous in the light distribution and perturb the circular velocityﬁeld through small-amplitude sinu-soidal disturbances(i.e.,wiggles in the rotation curves),but they are immaterial to the axisymmetric gravitational potential.This digression is to recall thatﬁt-ting these features with a mass model is a mistake!The rotation curves of spiral galaxies do not show any keplerian fall-oﬀat outer radii(e.g.:Rubin et al.1980;Bosma1981).Moreover,their shapes at R>(1−2)R D are inconsistent with the light distribution,so unveiling the presence of a DM component.PSS96,analyzing approximately1100RCs,about 100of which extended out to∼<2R opt,found that the luminosity speciﬁes the entire axisymmetric rotationﬁeld of spiral galaxies.At any chosen normalized4Figure3.Averaged spirals I-light proﬁles at diﬀerent luminosities.Each L bin includes hundreths of galaxiesradius x≡R/R opt,both the RC amplitude and the local slope strongly correlate with the galaxy luminosity(in particular,for x=1see Fig.2;for outer radii see PSS96,Salucci&Frenk1989,and Casertano&van Gorkom1991).Remarkably, the rms scatter around such relationships is much smaller than the variation of slopes among galaxies(see PSS96).This has led to the concept of the Universal Rotation Curve(URC)of spiral galaxies(PSS96and Persic&Salucci1991; see Fig.5).The rotation velocity of a galaxy of luminosity L/L∗at a radius x≡R/R opt is well described by:V URC(x)=V(R opt)0.72+0.44logLL∗1.97x1.22(x2+0.782)1.43+ +0.28−0.44logLL∗1+2.25LL∗0.4x2x2+2.25(LL∗)0.41/2km s−1.(3) (with log L∗/L =10.4in the B-band).Remarkably,spirals show a very small cosmic variance around the URC.In80%of the cases the diﬀerence between the individual RCs and the URC is smaller than the observational errors,and in most of the remaining cases it is due to a bulge not considered in eq.(3) (Hendry et al.1997;PSS96).This result has been conﬁrmed by a Principal Component Analysis study of URC(Rhee1996;Rhee&van Albada1997): they found that the twoﬁrst components alone account for∼90%of the total variance of the RC shapes.Thus,spirals sweep a narrow locus in the RC-proﬁle/amplitude/luminosity space.The luminosity dependence of the URC strongly contrasts with the self-similarity of the luminosity distribution of stellar disks(Fig.3):the luminosity proﬁles L(x)∝xx I(x)dx do not depend on luminosity.This reﬂects the5Figure4.Coadded rotation curves(ﬁlled circles with error bars) repruduced by URC(solid line)Also shown the separate dark/luminous contributions(dotted line:disk;dashed line:halo.)150300Figure5.The URC surface.6discrepancy between the distribution of light and that of the gravitating mass. Noticeably,this discrepancy increases with radius x and with decreasing galaxy luminosity L.The URC can beﬁtted by a combination of two components:(a) an exponential thin disk,approximated for0.04R opt<R≤2R opt asV2d(x)=V2(R opt)β1.97x1.22(x+0.78),(4)and(b)a spherical halo represented byV2h(x)=V2(R opt)(1−β)(1+a2)x2(x2+a2),(5)M h(x)=G−1V2(1)R opt(1−β)(1+a2)x3(x+a),with x≡R/R opt being the normalized galactocentric radius,β≡V2d(R opt)/V2opt the disk mass fraction at R opt,and a the halo core radius(in units of R opt).The disk+haloﬁts to the URC are extremely good(ﬁtting errors are within1%on average)at all luminosities(see Fig.4)whenβ=0.72+0.44logLL∗,(6)a=1.5LL∗1/5.(7)Thus we detect,for the DM component,a central constant-density region of size∼R opt,slightly increasing with luminosity.The transition between the inner,luminous-matter-dominated regime and the outer,DM-dominated regime occurs well inside the optical radius:typically at r<<R opt in low-luminosity galaxies,and farther out,closer to∼R opt,at high luminosities(see Fig.4).Thus, the ordinary and dark matter are well mixed in the very stellar regions of spirals.The total halo mass can be evaluated by extrapolating the halo out to the radius,R200,encompassing a mean overdensity of200.Weﬁnd:R200=250LL∗0.2kpc(8)M200∼2×1012LL∗0.5M ,(9)in good agreement with results derived from satellite and pair kinematics(Charl-ton&Salpeter1991;Zaritsky et al.1993).Notice thatM200 L B 75L BL∗−0.5(10)(in solar units).This implies that the halo mass function is not parallel to the observed galaxy luminosity function(Ashman,Salucci&Persic1993).7The luminosity dependence of the disk mass fraction can be interpreted in terms of a mass-dependent eﬃciency in transforming the primordial gas fraction into stars.In fact,the total(i.e.,computed at R200)luminous mass fraction of a galaxy of luminosity L is:M M200 0.05LL∗0.8.(11)This suggests that only the brightest objects reach a value comparable with the primordial valueΩBBN∼<0.10,while in low-L galaxies only a small fraction of their original baryon content has been turned into stars.The luminosity dependence of the DM fraction found by Persic&Salucci(1988,1990)has been conﬁrmed by direct mass modelling(e.g.:Broeils1992;Broeils&Courteau1997; Sincotte&Carignan1997;see also Ashman1992).On scales(0.2−1)R200,the halos have mostly the same structure,with a density proﬁle very similar at all masses and an amplitude that scales onlyvery weakly with mass,like V200∝M0.15200.This explains why the kinematics ofsatellite galaxies orbiting around spirals show velocities(relative to the primary) uncorrelated with the primary’s luminosity(Zaritsky1997).At very inner radii, however,the self-similarity of the proﬁle breaks down,the core radius becoming smaller for decreasing M200according to:core radiusR200=0.075M20010M0.6.(12)The central density scales with mass as:ρh(0)=6.3×104ρcM2001012M−1.3(13)(ρc is the critical density of the universe).The regularities of the luminous-to-dark mass structure can be represented as a curve in the space deﬁned by dark-to-luminous mass ratio,(halo core radius)-to-(optical radius)ratio,central halo density(see PSS96).This curve is a structural counterpart of the URC and represents the only locus of the mani-fold where spiral galaxies can be found.The main properties of dark matter in spiral galaxies can then be summarized as follows:substantial amounts of dark matter are detected in the optical regions of all spirals,starting at smaller radii for lower luminosities;the dark and the luminous matter are well mixed;the structure of the halos is universal:it involves a core radius compara-ble in size with the optical radius and a central density scaling inversely with luminosity;the ratio between luminous and dark matter is a function of luminosity(or mass)among galaxies.At a given(normalized)radius,this ratio increses with increasing luminosity:the global visible-to-dark mass ratio spans the range be-tween∼ΩBBN at very high luminosities and∼10−4at L<<L∗.The discovery of these features,describing the various stages of the pro-cesses leading to present-day galaxies,supersedes the observationally disproved paradigms of”ﬂat rotation curves”and”cosmic conspiracy”.8Figure 6.Coadded rotation curve of the sample of LSB galaxies with V opt ∼(70±30)km s −1.The solid line represents the V URC of spirals of similar V opt .3.Low-Surface-Brightness GalaxiesThe central surface luminosity of spirals is normally about constant,at µ0(B )=21.65±0.30.Recently,however,a large population of disk systems with a signif-icantly lower surface brigthness (µ0(B )=24−25)has been detected (Schombert et al.1992;Driver et al.1994;Morshidi et al.1997;see also Disney &Phillips 1983).In these systems the light distribution follows that of an exponential thin disk (McGough &Bothun 1994),with total magnitude ∼1.5mag fainter than that of normal (HSB)spirals.At the faint end of the LSB luminosity distribu-tion (M B ∼−16),the galaxies have very extended HI disks whose kinematics yields the mass structure (de Blok,McGaugh &van der Hulst 1996).In detail,de Blok et al.have published the HI rotation curves of 19low-L LSBs,having M B ∼−17mag,V max ∼70km s −1,and (remarkably)R opt ∼(8−10)kpc,roughly independent of luminosity.These galaxies are the counterpart of the faintest HSB spirals.Notice that,although the maximum velocities are similar,the optical sizes of LSBs are ∼3times larger than those of HSBs.The objects in de Blok et al.(1996)are all within a small range of magni-tudes and optical velocities,M B =−17±0.5,V opt ∼(70±30)km s −1.From these data we construct,as we did for spirals,the coadded rotation curve V (R R opt ,70).Notice that,unlike for spirals,rotation curves of LSBs with higher V opt are not available:so we can compare LSB and spiral RCs only at their lowest velocities,shown for spirals in the top left panel in Fig.4and for LSBs in Fig.6(points).In detail,in the latter ﬁgure we plot V (R/R opt ),the coadded RC obtained from the de Blok et al.data,together with the spiral V URC for the same range of maximum velocities (50–100km s −1)(see PSS96for all of the details).The9Figure 7.The radius–luminosity relation (left)and the Tully-Fisher relation (right)for the Matthews sample of LSBs.Dashed curve:ﬁt to the data.Solid curve:predictions of a mass model identical to that of spirals.Dotted line:prediction of a mass model model with constant (M/L )agreement is striking:no diﬀerence can be detected between the LSB and HSB rotation curves,which coincide within the observational uncertainties.Notice that the LSB and HSB rotation curves are identical when the radial coordinate is normalized to the disk length-scale (essential procedure for determining the DM distribution).The fact that,when expressed in physical radii,LSB rotation curves rise to their maximum more gently than HSB RCs (see Fig.3of de Blok &McGough 1997),depends exclusively on the larger LSB disk scalelengths.For LSBs we adopt a mass model including,as for spirals,(i)an exponential thin disk and (ii)a dark halo of mass proﬁle given by eq.(5).The LSB coadded RC is extremely well ﬁtted by eqs.(4)and (5)with the ‘spiral’values β=0.1±0.04and a 0.75.These parameters imply that LSB galaxies are completely dominated by a dark halo with a large ∼(5−6)kpc core radius.Higher-luminosity (up to M B ∼−22)LSBs do exist (e.g.,Sprayberry et al.1993):their structure can be tentatively investigated by means of their Tully-Fisher relationship.Observationally,these objects show a good correlation between luminosity and the corrected linewidth w 0,i (Matthews et al.1997):log w 0,i =2.41+0.69log L V L ∗V +0.16log 2 L V L ∗V ,(14)where L ∗V corresponds to M ∗V =−20.87(see Fig.7).The quantity w 0,i is,as inspirals,a good measure of the circular velocity at ∼R opt :thenw 20,i R opt G [M h (R opt )+M bar (R opt )],(15)where M bar =M D +M HI .For LSB spirals we take the same dark-to-visible mass ratios as for HSB spirals,M h (R opt )M bar (R opt ) 9×(L B 0.04L ∗)−1(see conclusions of PSS96),and we assume M bar (R opt )∝L k B .Finally,by using the LSB luminosity-radius relation,R opt =12(L/L ∗)0.25kpc (see Fig.7),we are able to reproduce10Figure8.LSB(M/L) (left:solid curve;the dotted line refers tospirals)and central overdensity(right)as a function of luminosity.the observed M V–log w0,i relation:in Fig.7we show the excellent agreement between the observed linewidths and those predicted by our mass model when k=1.8and the LSBs stay on the‘spiral’β–L and a–L relationships.At∼L∗,the LSB stellar M/L ratios are similar to those of HSBs,but they decline steeply with decreasing luminosity(see Fig.8,left panel).This result is in good agreement with the(M/L) ratios obtained by applying the maximum-disk hypothesis to solve for the galactic structure:in this case log(M/L) increases from∼−0.3to0.5across the entire LSB luminosity(velocity)range(de Blok &McGaugh1997).The central overdensities,3/(4πGρc)(1−β)(V opt/R opt)2(1+a2)/a2,are shown in Fig.8(right panel)as a function of luminosity.They are smaller and less dependent on galaxy luminosity than those in normal spirals,in agreement with theﬁndings by de Blok&McGough(1997).With the caveat of the relative smallness of the present sample,we claim that LSB galaxies are indistinguishable from HSBs in the(β,a)space describing the coupling between dark and visible matter.To end this section,we comment on the argument,raised sometimes,accord-ing to which the discovery of a large number of LSBs would make theﬁndings on DM in spirals(e.g.PSS96)irrelevant,in that normal spirals would represent an unfair,biased sample of galaxies.There is no doubt that the existence of a population of LSB galaxies,contiguous to the HSBs,aﬀects the interpretation, and even the physical meanings,of luminosity functions and number counts. However,it is exactly the characteristic diﬀerence between the two families that answers the point.Even prior to any mass modelling,we can argue that in the Universe there are about1011disk systems obeying the“Freeman law”:µ0=21.5±0.5independent of galaxy luminosity.It is obviously necessary to know the DM properties of this family,independently of the existence of another family of disk systems not obeying such a“law”,and maybe not even observ-able.Of course,investigating the DM properties of LSBs is equally important and crucial.Furthermore,the argument raised above appears even more im-material after the analysis of the LSB rotation curves.In fact,as far as the structural parameters are concerned,HSBs and LSBs are indistinguishable.In11Figure9.The mass as a function of radius for the ellipticals of oursubsamples,with luminosities<log L/L∗>=−0.4and<log L/L∗=0.4(left and right,respectively).other words,LSBs follow the spiral V URC(R/R opt),but strongly deviate with respect to the spirals’luminosity–(disk length-scale)relation.4.Elliptical GalaxiesElliptical galaxies are pressure-supported,triaxial stellar systems whose orbital structure may depend on their angular momentum content,degrees of triaxial-ity and velocity dispersion anisotropy.As is well known,the derivation of the mass distribution from stellar motions is not straightforward as it is from ro-tation curves in disk systems,because the kinematics of the former is strongly aﬀected by geometry,rotation,and anisotropies.However,in addition to stellar kinematics,there are a number of mass tracers(X-ray emitting halos,planetary nebulae,ionized and neutral disks)which crucially help to probe the gravita-tional potential out to external radii(see Danziger1997).The luminosity proﬁleρ (r)of E’s is obtained from the observed projected surface density,which follows the de Vaucouleurs proﬁle[see eq.(1)],by assuming an intrinsic shape and deprojecting.A very good approximation forρ (r)is theHernquist proﬁle:ρ (r)=M2π1y(y+c)(16)(where y=1.81r/r e and c=1).As the Hernquist proﬁle is no longer a good ﬁt for R>>r e,from actual surface-photometry proﬁles we have evaluated thatR opt 2r e,where r e=6(L/L∗)0.7kpc(from Djorgovski&Davis1987).All ellipticals belong,within a very small cosmic scatter(<12%),to a relation ofthe type:r e∝σA0I B e,(17) where r e is the half-light radius,I e is the mean surface brightness within r e,and σ0is the observed(projected)central velocity dispersion.In the logarithmic space,this corresponds to a plane,the Fundamental Plane of ellipticals(Djor-1200.51 1.52012300.511.520123Figure 10.Coadded mass proﬁle of ellipticals,ﬁlled cirles,with thebest ﬁt mass model (solid line)govski &Davis 1987;Dressler et al.1987),that constrains the properties of the DM distribution.Observations indicate A =1.23(1.66)and B =−0.82(−0.75)in the V -band and K -band,respectively (e.g.,Djorgovski &Santiago 1993).Assuming a constant M/L and structural homology,the virial theorem predicts A =2,B =−1.The simplest explanation of the departure of A from the virial expectation involves a systematic variation of M/L with L ,which also accounts for the wavelength dependence of A (Djorgovski &Santiago 1993).The departure of B from the virial expectation is likely to be due to the breakdown of the homology of the luminosity structure (see Caon,Capaccioli &D’Onofrio 1993;Graham &Colless 1997).Assuming spherical symmetry and isotropic stellar motions (so M 3.4G −1σ20r e and L =2πI e r 2e ),the FP implies that M/L |r e ,inclusive of dark matter,is “low”:4−9(Lanzoni 1994;Bender,Burstein &Faber 1992),and roughly consistent with the values predicted by the stellar population models (Tinsley 1981).No large amounts of DM are needed on these scales,as it also emerges from mass models of individual ellipticals,obtained by analyzing the line proﬁles of the l.o.s velocity dispersion (van der Marel &Franx 1993),which show that the DM fraction inside r e is substantially less than 50%(i.e.M/L B ∼<10;see van der Marel 1991,1994;Rix et al.1997;Saglia et al.1997a,b;Carollo &Danziger 1994;Carollo et al.1995),as in spirals.We now investigate in some detail the eﬀects of DM on the Fundamental Plane.For this purpose,let us describe,for mathematical simplicity,the dark halo by a Hernquist proﬁle [eq.(16)]with a lower mass concentation than the luminous spheroid:c =2(see Lanzoni 1994).We recall that dark halos are likely to have an innermost constant-density region not described by eq.(16):this,13however,has no great relevance here,in that most of the dark mass is located outside the core where eq.(16)is likely to hold.Without loss of generality,we consider isotropic models (Lanzoni 1994):the radial dispersion velocity σr is related,through the Jeans equation,to the mass distribution by:σ2r (r )=G M dark (r )+M (r )r d log ρd log r −1.(18)The projected velocity dispersion is then σP (R )=2Σ (R ) ∞R ρ (r )σ2r (r )√2−R 2dr .This equation shows that,at any radius (including r =0),the measured projected ve-locity dispersion σP (R )depends on the distributions of both dark and luminous matter out to r ∼(2−3)r e .Conversely,in spirals the circular velocity V (R )depends essentially only on the mass inside R ,namely on just the luminous mass when R →0.If M 200is the total galaxy mass,we getσP (0)=3.4GM r e M 20010M 2/5.(19)The mass dependence of σP (0),combined with the thinness of the FP (i.e.:δR e /R e ∼<0.12),constrains the scatter that would arise,according to eq.(19),from random variations of the total amount of DM mass in galaxies with the same luminous mass.From eq.(17)we get δσ/σ∼<0.121.4,while eq.(19)implies dσ/σ=5/2δM/M .This means that,over 2orders of magnitude in M ,any random variation of the dark mass must be less than 20%(Lanzoni 1994;Renzini &Ciotti 1993;Djorgovski &Davis 1987).From eq.(17)and since M ∝L 1.2,we ﬁnally get M M 200∝σ5/2,so that M 200∝L 0.5−0.6,as in spirals (PSS96).Let us notice that the these well-established constraints on the amount of darkmatter in ellipticals,due to the existence of the Fundamental Plane,are however at strong variance with the extremely high values of the central M/L ratios,15−30,found in some objects,as a result of the dynamical models of Bertin et al.(1992)and Danziger (1997).We can determine the parameters of the dark and visible matter distribu-tion by means of a variety of tracers of the ellipticals’gravitational ﬁeld (see the review by Danziger 1997).This provides the (dark +luminous)mass dis-tribution for 12galaxies,with magnitudes ranging between −20<M B <−23.In order to investigate the luminosity dependence of the mass distribution,we divide the sample into 2subsamples,with average values <log L/L ∗>=−0.4and <log L/L ∗>=0.4respectively.In Fig.9we plot,as a function of R/r e ,the normalized mass proﬁle of ellipticals,M (R )/M (r e ),for each galaxy;and in Fig.10the coadded mass distributions for the high-L and the low-L subsamples that can be very well reproduced (solid lines)by a two-component mass model which includes a luminous Hernquist spheroid and a DM halo given by eq.(5).The resulting ﬁt is excellent (see Fig.10)when β,the luminous mass fraction inside R opt 2r e ,scales with luminosity as in disk systems,and the halo core radius a (expressed in units of R opt )scales asa =0.8L L ∗ 0.15.(20)14。

1993年2月GRE考试真题

1993年2月GRE考试真题SECTION 31. The corporation expects only ------- increases in sales next year despite a yearlong effort to revive itsretailing business.(A) unquestionable(B) sequential(C) modest(D) exaggerated(E) groundless2. No computer system is immune to a virus, a particularly malicious program that is designed to -------and electronically -------- the disks on which data are stored.(A) prepare.. improve(B) restore.. disable(C) infect.. damage(D) preserve.. secure(E) invade.. repair3. Recent research indicates that a system of particles which has apparently decayed to randomness from------- state can be returned to that state; thus the system exhibits a kind of memory of its ------- condition.(A) an equilibrium.. lesser(B) an ordered.. earlier(C) an unusual.. settled(D) a chaotic.. last(E) a higher.. present4. A number of writers who once greatly ------- the literary critic have recently recanted, substituting------- for their former criticism.(A) lauded.. censure(B) influenced.. analysis(A) simulated.. ambivalence(B) disparaged.. approbation(C) honored.. adulation5. She writes across generational lines, making the past so ------- that our belief that the present is the true locus of experience is undermined.(A) complex(B) distant(C) vivid(D) mysterious(E) mundane6. Individual freedom of thought should be ------- more absolutely than individual freedom of action, giventhat the latter, though also desirable, must be ------- the limits imposed by the rights and freedom ofothers.(A) protected.. subject to(B) assessed.. measured by(C) valued.. superior to(D) exercised.. indifferent to(E) curtailed.. conscious of7. Their ----- was expressed in quotidian behavior: they worshipped regularly, ---- all the regenerative processed of nature respect, and even awe.(A) selflessness.. reserving to(B) moderation.. extending to(C) reverence.. exacting from(D) piety.. according(F) serenity.. refusing8. CHAFF: WHEAT::(A)spore: seed(B)nucleus : cell(C)sod : flower(D)shell : pecan(E)root : tooth9. ARRAY : NUMBERS(A)body : skeleton(B)formation : soldiers(C)club: members(D)rank: insignia(E)illustration: graphs10. MASK : FACE::(A)pseudonym: name(B)caricature : likeness(C)forgery: imitation(D)disguise: detective(E)code: agent11. INCORRIGIBLE : REFORMED::(A)inscrutable : understood(B)infallible : believed(C)inferior : defeated(D)ingenious: copied(E)infamous : condemned12. FILIBUSTER : LEGISLATION::(A)restriction: zone(B)blockade: commerce(C)suspension : sentence(D)denial : accusation(E)intermission: performance13. FROND : LEAF::(A)lawn: grass(B)wasteland : water(C)thicket: shrub(D)river: pond(E)boulder: rock14. TINT: SUFFUSE::(A)ponder: yearn(B)regret: undo(C)damp: quench(D)shroud : screen(E)amble: wander15. MAGAZINE: PERIODICAL::(A)newspaper: edition(B)mystery : fiction(C)volume: encyclopedia(D)chapter: book(E)article: journal16. FRANK: SECRETIVENESS::(A)honest: theft(B)transparent : light(C)free: autocracy(D)callow: maturity(E)confident : intrepidnessIt is now established that the MilkyWay is far more extended and of muchgreater mass than was hitherto thought.However, all that is visible of the(5) constituents of the Milky Way's corona(outer edge), where much of the galaxy's mass must be located, is a tiny fraction of the corona's mass. Thus, most of theMilky Way's outlying matter must(10) be dark.Why? Three facts are salient. First,dwarf galaxies and globular clusters,into which most of the stars of theMilky Way's corona arc probably bound,(15) consist mainly of old stars. Second,old stars are not highly luminous.Third, no one has detected in the coronathe clouds of gaseous matter such ashydrogen and carbon monoxide that are(20) characteristic of the bright parts of agalaxy. At present, therefore, the bestexplanation-though still quitetentative--for the darkness of thecorona is that the corona is composed(25) mainly of old, burned-out stars.17. The passage as a whole is primarily concerned with(A) analyzing a current debate(B) criticizing a well-established theory(C) showing how new facts support a previously dismissed hypothesis(D) stating a conclusion and adducing evidence that may justify it(E) contrasting two types of phenomena and showing how they are related18. According to the passage, a bright part of a galaxy typically includes(A) dwarf galaxies and clusters of stars(B) a balanced mixture of old and new stars(C) a large portion of the galaxy’s mass(D) part of the corona of the galaxy(E) gases such as hydrogen and carbon monoxide19. It can be inferred from the passage that, compared with what they now think,until fairly recently astronomers believed that the Milky Way(A) was much darker(B) was much smaller(C) was moving much more slowly(D) had a much larger corona(E) had much less gaseous matter20. The passage presents which of the following as incontrovertible?I. The low luminosity of old starsII. The absence of clouds of gaseous matter from the corona of the Milky WayIII. The predominance of globular clusters and dwarf galaxies in the corona of the Milky Way(A) I only(B) III only(C) I and II only(D) II and III only(E) I, II, and IIIOne of the principal of Walzer'scritique of liberal capitalism isthat it is insufficiently egali-tarian. Walzer's case against the(5) economic inequality generated bycapitalism and in favor of "aradical redistribution of wealth"is presented in a widely citedessay entitled "In Defense of(10) Equality."The most striking feature ofWaizer's critique is that, farfrom rejecting the principle ofreward according to merit, Walzer(15) insists, on its validity. Peoplewho excel should receive thesuperior benefits appropriate totheir excellence. But people exhibita great variety of qualities-"intelli-(20) gence, physical strength. agility andgrace. artistic creativity, mechanicalskill. leadership, endurance, memory,psychological insight. the capacityfor hard work-even moral strength, (25) sensitivity. the ability to expresscompassion." Each deserves its properrecompense. and hence a proper distri-bution of material goods should reflecthuman differences as measured on all (30) these different scales. Yet, undercapitalism, the ability to make money("the green thumb of bourgeois society") enables its possessor to acquire almost"every other sort of social good." (35) such as the respect and esteem ofothers.The centerpiece of Walzer's argument is the invocation of a quotation fromPascal's Pensees, which concludes:(40) "Tyranny is the wish to obtain byone means what can only be had byanother." Pascal believes that weowe different duties to differentqualities. So we might say that(45) infatuation is the proper responseto charm, and awe the proper responseto strength. In this light, Walzercharacterizes capitalism as thetyranny of money (or of the ability(50) to make it) And Walzer advocates asthe means of eliminating this tyrannyand of restoring genuine equality"the abolition of the power of moneyoutside its sphere" What Walzer envi- (55) sions is a society in which wealth isno longer convertible into social goodswith which it has no intrinsic connection.Walzer's argument is a puzzling one.After all, why should those qualities (60) unrelated to the production of materialgoods be rewarded with material goods?Is it not tyrannical, in Pascal's sense,to insist that those who excel in "sensi-tivity" or "the ability to express(65) compassion" merit equal wealth withthose who excel in qualities (such as"the capacity for hard work") essentialin producing wealth? Yet Waizer'sargument, however deficient, does(70) point to one of the most serious weak-nesses of capitalism-namely, that itbrings to predominant positions in asociety people who, no matter howlegitimately they have earned their(75) material rewards, often lack thoseother qualities that evoke affectionor admiration. Some even argue plausiblythat this weakness may be irremediable:in any society that, like a capitalist(80) society, seeks to become ever wealthierin material terms disproportionaterewards are bound to flow to the peoplewho are instrumental in producing theincrease in its wealth.21. The primary purpose of the passage is to(A) argue that Walzer’s critique of liberal capitalism is the cornerstoneof Walzer’s thinking(B) identify and to deprecate the origins of the intellectual traditionchampioned by Walzer(C) present more clearly than does the essay “In Defense of Equality”thedistinctive features of Walzer’s politico-economic theories(D) demonstrate that Walzer’s critique of liberal capitalism is neitheroriginal nor persuasive(E) outline and to examine critically Walzer’s position on economicequality22. The author mentions all of the following as issues addressed by WalzerEXCEPT:(A) proper recompense for individual excellence(B) proper interpretation of “economic equality”(C) proper level of a society’s wealth(D) grounds for calling capitalism “the tyranny of money”(E) exchangeability of money for social goods23. The argumentation in the passage turns importantly on the question of whatshould be the proper relation between(A) “liberal capitalism” (line 2) and “bourgeois society” (lines 20-21)(B) “reward” (line 8) and “recompense” (line 17)(C) “sensitivity”(line 15) and “the ability to express compassion”(lines15-16)(D) “distribution of material goods”(lines 17-18) and “redistribution ofwealth” (lines 4-5)(E) “social goods” (line 37) and “material goods” (line 41)24. The passage provides sufficient information to answer which of thefollowing questions?(A) What weight in relation to other qualities should a quality likesensitivity have, according to Walzer, in determining the properdistribution of goods?(B) Which quality does Walzer deem too highly valued under liberalcapitalism?(C) Which are the social goods that are, according to Walzer, outside thereach of the power of money?(D) What practical steps does Walzer suggest be taken to relieve theeconomic inequality generated by capitalism?(E) What deficiencies in Walzer’s own argument does Walzer acknowledge?25. The author implies that Walzer’s interpretation of the principle of rewardaccording to merit is distinctive for its(A) insistence on maximizing everyone’s rewards(B) emphasis on equality(C) proven validity(D) broad conception of what constitutes merit(E) broad conception of what constitutes a reward26. The author’s interpretation of the principle that “we owe different dutiesto different qualities”(lines 28-29) suggests that which of the following would most probably be the duty paired with the quality of veracity?(A) Dignity(B) Trust(C) Affection(D) Obedience(E) Integrity27. The author implies that sensitivity is not a quality that(A) is essential in producing wealth(B) wealthy people lack(C) can be sensibly measured on a scale(D) characterizes tyrannical people(E) is owed a duty in Pascal’s sense28. SYMMETRY:(A)separateness(B)corruption(C)mutability(D)imprecision(E)disproportion29. DIVERGENCE:(A)peacefulness(B)control(C)stipulation(D)contentment(E)unification30. OBSTRUCTIONIST:(A)one who governs(B)one who welcomes(C)one who repents(D)one who facilitates(E)one who trusts31. DIURNAL:(A)nomadic(B)aggressive(C)cold-blooded(D)chiefly active at night(E)often randomly distributed32. AXIOMATIC:(A)controversial(B)peremptory(C)uncomplicated(D)vestigial(E)amalgamated33. SUBVERT:(A)increase(B)replace(C)reinforce(D)oversee(E)expose34. FOMENT:(A)simplify(B)rectify(C)isolate(D)explain(E)stifle35. ENNUI:(A)annoyance(B)excitement(C)sympathy(D)misery(E)assurance36. EQUABLE:(A)boundless(B)intemperate(C)tangential(D)flimsy(E)pernicious37. HUBRIS:(A)mockery(B)calm(C)confusion(D)approval(E)humility38. SURFEIT:(A)select(B)caution(C)repose(D)starve(E)consoleSECTION 6Time-30 minutes38 Questions1. My family often found others laughable, but I learned quite early to be -------- while people werepresent, laughing only later at what was funny andmocking what to us seemed(A) polite.. bizarre(B) impatient.. unfortunate(C) facetious.. enviable(D) wistful.. extraordinary(E) superficial.. deplorable2. The technical know-how, if not the political ------- ,appears already at hand to feed the world's explodingpopulation and so to ----- at last the ancient scourges of malnutrition and famine.(A) will.. weaken(B) expertise.. articulate(C) doubt.. banish(D) power.. denounce(E) commitment.. eradicate3. In small farming communities, accident victims rarely sue or demand compensation: transforming a personal injury into a ------- someone else is viewed as an attempt to ------- responsibility for one's own actions.(A) conspiracy against.. assume(B) claim against.. elude(C) boon for. .minimize(D) distinction for.. shift(E) trauma for.. proclaim4. Dominant interests often benefit most from ------- of governmental interference in business, since they are able to take care of themselves if left alone.(A) intensification(B) authorization(C) centralization(D) improvisation(E) elimination5. The "impostor syndrome" often afflicts those whofear that true self-disclosure will lower them in others'esteem; rightly handled, however, ------- may actually------ one's standing.(A) willfulness.. consolidate(B) imposture.. undermine(C) affectation.. jeopardize(D) candor. .enhance(E) mimicry.. efface6. The pungent verbal give-and-take among thecharacters makes the novel ------ reading, and thisvery ------- suggests to me that some of the opinionsvoiced may be the author's.(A) disturbing.. flatness(A) tedious. inventiveness(B) lively.. spiritedness(D) necessary.. steadiness(E) rewarding.. frivolousness7. The fortresslike facade of the Museum of CartoonArt seems calculated to remind visitors that the comic strip is an art form that has often been——by critics.(A) charmed(B) assailed(C) unnoticed(D) exhilarated(E) overwhelmed(A)8. SPLICE: ROPE::(B)press :shirt(C)caulk: frame(D)weld: metal(E)plaster: wall(F)curl: hair9. FANATIC: DEVOTED::(A)prude: proper(B)skeptic: religious(C)cad: devious(D)gourmet : ravenous(E)coquette: graceful10. CONFLUENCE: STREAMS::(A)ridge: hills(B)railroad: tracks(C)junction: roads(D)curb: sidewalks(E)park: edges11. SWAGGER: BRAVADO::(A)chevron : sergeant(B)sword : bravery(C)salute : disrespect(D)caress : affection(E)sneeze: explosion12. INDECOROUS: PROPRIETY::(A)boorish : sensitivity(B)rancorous : hostility(C)stuffy: dignity(D)presumptuous: boldness(E)charismatic: loyalty13. CAPRICIOUS: WHIM::(A)conventional : innovation(B)objective : fact(C)satirical : benevolence(D)gloomy : optimism(E)opinionated : rudeness14. SNOW: PRECIPITATION::(A)lava: volcano(B)hurricane: cyclone(C)desert: drought(D)seed : germination(E)temperature : season15. RECALCITRANT: AUTHORITY::(A)implacable: conciliation(B)remorseful: recompense(C)indomitable: challenge(D)insubordinate: camaraderie(E)enthusiastic: opportunity16. INKLING: INDICATON::(A)apprentice: expert(B)theory: hypothesis(C)hunger: thirst(D)orientation: direction(E)lapse : errorThe outpouring of contemporaryAmerican Indian literature in the last two decades, often called the NativeAmerican Renaissance, represents for(5) many the first opportunity toexperience Native American poetry. Theappreciation of traditional oralAmerican Indian literature has beenlimited, hampered by poor translations (10) and by the difficulty, even in the rareculturally sensitive and aestheticallysatisfying translation, of completelyconveying the original's versestructure, tone, and syntax.(15) By writing in English and experimentingwith European literary forms,contemporary American Indian writershave broadened, their potentialaudience, while clearly retaining many(20) essential characteristics of theirancestral oral traditions. Forexample, Pulitzer-prize-winning authorN. Scott Momaday's poetry often treatsart and mortality in a manner that(25) recalls British romantic poetry, whilehis poetic response to the power ofnatural forces recalls Cherokee oralliterature. In the same way, hisnovels, an art form European in origin,(30) display an eloquence that echoes theoratorical grandeur of the greatnineteenth- century American Indianchiefs.17. According to the passage, Momaday’s poetry shares which of the followingwith British romantic poetry?(A) Verse structure(B) Oratorical techniques(C) Manner of treating certain themes(D) Use of certain syntactical constructions(E) Patterns of rhythm and rhyme18. Which of the following is most likely one of the reasons that the authormentions the work of N. Scott Momaday?(A) To illustrate how the author believes that members of the NativeAmerican Renaissance have broadened their potential audience(B) To emphasize the similarities between Momaday’s writings and theirEuropean literary models(C) To demonstrate the contemporary appeal of traditional Native Americanoral literature(D) To suggest that contemporary American Indian writers have sacrificedtraditional values for popular literary success(E) To imply the continuing popularity of translations of oral AmericanIndian literature19. Which of the following can be inferred from the passage about writtentranslations of oral Native American poetry?(A) They were less widely read than are the works of contemporary NativeAmerican poets writing in English.(B) They were often made by writers who were intimately familiar with bothEnglish and Native American languages.(C) They often gave their readers aesthetic satisfaction, despite theirinaccuracies.(D) They usually lacked complex verse structure.(E) They were overly dependent on European literary models.20. The passage suggests which of the following about American Indian poetsbefore the Native American Renaissance?(A) Art and mortality were rarely the subjects of their poetry.(B) Their oratorical grandeur reached its peak in the nineteenth century.(C) They occasionally translated their own poetry.(D) They seldom wrote poetry in English.(E) They emphasized structure, tone, and syntax rather than literary form.Recent findings suggest that visualsignals are fed into at least threeseparate processing systems in thebrain, each with its own distinct(5) function. One system appears, toprocess information about shapeperception; a second, informationabout color; a third, informationabout movement, location, and spatial(10) organization. An understanding ofthe functions and capabilities ofthese three systems can shed lighton how artists manipulate materialsto create surprising visual effects.(15) It is possible to summarize thefunctions of the three subsystemsof the visual system as follows.The parvo system carries highlydetailed information about stationary(20) objects and about borders that areformed by contrasting colors. Itdoes not, however, carry informationabout specific colors. Because muchof the information about the shape of(25) objects can be represented by theirborders, we suspect that this systemis important in shape perception. Theblob system processes information about colors, but not about movement, shape (30) discrimination, or depth. The magnosystem carries information aboutmovement and depth. It is good atdetecting motion but poor at scru-tinizing stationary images. In(35) addition it appears to be colorblind;it is unable to perceive borders thatare visible only on the basis of colorcontrast.Cells in parvo system can distinguish (40) between two colors at any relativebrightness of the two. Cells in thecolor-blind magno system. on the otherhand, are analogous to a black-and-white photograph in the way they(45) function: they signal informationabout the brightness of surfaces butnot about their colors. For any pairof colors there is a particular bright- ness ratio at which two colors, for (50) example red and green, will appear asthe same shade of gray in a black-and-white photograph, hence any borderbetween them will vanish. Similarlyat some relative red-to-green bright- (55) ness level, the red and green willappear identical to the magno system.The red and green are then called equi-luminant. A border between two equi-luminant colors has color contrast(60) but no luminance contrast.Many artists have seemed to beempirically aware of these underlyingprinciples and have used them tomaximize particular effects. Some(65) of the peculiar effects of Op Art,for example, probably arise fromcolor combinations that are strongactivators of the parvo system butare weak stimuli for the magno system.(70) An object that is equiluminant withits background looks vibrant andunstable. The reason is that theparvo system can signal the object'sshape but the magno system cannot see(75) its borders and therefore cannot signaleither the movement or the position ofthe object. Hence it seems to jumparound, drift, or vibrate on the canvas.21. The passage is primarily concerned with(A) describing subsystems of the visual system and showing their relevanceto art(B) comparing three theories on how the visual system analyzes images ina work of art(C) explaining how artists use color contrasts to create particular visualeffects(D) explaining how the visual system distinguishes among different colors(E) describing functions of the first three phases of the visual system22. Which of the following would create visual effects most similar to thosediscussed in lines 43-48?(A) A watercolor in which colors are applied imprecisely to outlined shapes(B) A painting in which different shades of the same color are used toobscure the boundaries between objects(C) A black-and-white sketch in which shading is used to convey a senseof depth(D) An advertisement in which key words are at the same level of brightnessas a background of contrasting color(E) A design in which two different shades of gray are juxtaposed toheighten the contrast between them23. The passage provides information about which of the following?(A) Why the same system can process information about movement and location(B) Why the parvo system is considered to be responsible for shapeperception(C) Why the blob system can process information about colors but notmovement(D) The mechanism that enables the blob system to distinguish betweenstationary objects(E) The mechanism that enables the magno system to carry information aboutshape discrimination24. According to the passage, which of the following is true of the visualsystem?(A) It processes visual signals in three consecutive stages.(B) It processes visual signals through separate processing systems inthe brain.(C) It consists of only three separate systems.(D) It consists of a single hierarchical system rather than a multipartitesystem.(E) It consists of separate system with high overlap in processingfunctions.25. The author mentions a “black-and-white photograph”(line 29) most probablyin order to explain(A) how the parvo system distinguishes between different shapes and colors(B) how the magno system uses luminosity to identify borders betweenobjects(C) the mechanism that makes the magno system color-blind(D) why the magno system is capable of perceiving moving images(E) the brightness ratio at which colors become indistinguishable to theparvo system26. The author uses all of the following in the discussion in the thirdparagraph EXCEPT:(A) an example(B) definition of terms(C) contrast(D) a rhetorical question(E) analogy27. The passage suggests which of the following about the magno system?(A) It perceives borders on the basis of luminance contrast.(B) It perceives shapes on the basis of color contrast.(C) It is better at perceiving stationary objects than it is at detectingmovement.(D) It can detect motion but it cannot signal the position of an object.(E) It is better at processing information about movement than it is atprocessing information about depth.28. MODISH:(A)eliciting admiration and joy(B)avoiding harm and danger(C)lacking style and fashionableness(D)providing vitality and fortitude(E)destroying usefulness and serviceability29. SPINY:(A)heavy(B)placid(C)smooth(D)terse(E)single30. SCRUTINIZE(A)demur(B)dispute(C)condone(D)elaborate on(E)gloss over31. INCLEMENT:(A)torpid(B)truculent(C)buoyant(D)balmy(E)bucolic32. RAZE:(A)build(B)strengthen(C)impede(D)refurbish(E)stabilize33. PANDEMIC:(A)unaware(B)disapproving(C)soothing(D)faultless(E)limited34. EXCORIATE:(A)accept conditionally(B)praise lavishly(C)esteem grudgingly(D)permit(E)relax35. GILD:(A)prepare carelessly(B)offer hesitantly(C)represent accurately(D)speak forcibly(E)organize coherently36. RAREFY:(A)concentrate(B)modulate(C)diversify(D)leave(E)waste37. ASPERSION:(A)mandate(B)covenant(C)heartfelt gratitude(D)solemn declaration(E)glowing tribute38. PERSPICUITY(A)opacity(B)unrelatedness(C)fragility(D)unfamiliarity(E)deviance。

全球首个促睫毛生长处方药Latisse(比马前列素眼科溶液

HIGHLIGHTS OF PRESCRIBING INFORMATIONThese highlights do not include all the information needed to use LATISSE® safely and effectively. See full prescribing information for LATISSE®.LATISSE® (bimatoprost ophthalmic solution) 0.03%Initial U.S. Approval: 2001____________________INDICATIONS AND USAGE_____________________ LATISSE® is a prostaglandin analog, indicated to treat hypotrichosis of the eyelashes by increasing their growth including length, thickness and darkness. (1)________________DOSAGE AND ADMINISTRATION________________ Apply nightly directly to the skin of the upper eyelid margin at the base of the eyelashes using the accompanying applicators. Blot any excess solution beyond the eyelid margin. Dispose of the applicator after one use. Repeat for the opposite eyelid margin using a new sterile applicator. (2)______________DOSAGE FORMS AND STRENGTHS_______________ Bimatoprost ophthalmic solution 0.3 mg/mL. (3) ________________________CONTRAINDICATIONS_______________________ Hypersensitivity. (4.1)_________________WARNINGS AND PRECAUTIONS________________ Concurrent administration of LATISSE® and IOP-lowering prostaglandin analogs in ocular hypertensive patients may decrease the IOP-lowering effect. Patients using these products concomitantly should be closely monitored for changes to their intraocular pressure. (5.1)Pigmentation of the eyelids and iris may occur. Iris pigmentation is likely to be permanent. (5.2, 5.3)_______________________ADVERSE REACTIONS________________________ Most common adverse events (incidence approximately 3% - 4%) are eye pruritus, conjunctival hyperemia, and skin hyperpigmentation. (6)To report SUSPECTED ADVERSE REACTIONS, contact Allergan at 1-800-433-8871 or the FDA at 1-800-FDA-1088 or /medwatch.See 17 for Patient Counseling Information and FDA-approved Patient Package Insert.Revised: 07/2009FULL PRESCRIBING INFORMATION: Contents* 1 INDICATIONSANDUSAGE2 DOSAGEANDADMINISTRATION3 DOSAGE FORMS AND STRENGTHS4 CONTRAINDICATIONS4.1Hypersensitivity5 WARNINGSANDPRECAUTIONS5.1 Effects on Intraocular Pressure5.2IrisPigmentation5.3 Lid Pigmentation5.4 Hair Growth Outside the TreatmentArea5.5IntraocularInflammation5.6MacularEdema5.7ContaminationofLATISSE® orApplicators5.8 Use with Contact Lenses6 ADVERSEREACTIONS8 USEINSPECIFIC POPULATIONS8.1Pregnancy8.3NursingMothers8.4PediatricUse8.5GeriatricUse11 DESCRIPTION12 CLINICALPHARMACOLOGY12.1 MechanismofAction12.3 Pharmacokinetics13 NONCLINICALTOXICOLOGY13.1 Carcinogenesis,Mutagenesis,ImpairmentofFertility14 CLINICALSTUDIES16 HOW SUPPLIED/STORAGE AND HANDLING17 PATIENT COUNSELINGINFORMATION17.1 NightlyApplication17.2 Handling the Bottle and Applicator17.3 Potential for Intraocular PressureEffects17.4 Potential for Eyelid Skin Darkening17.5 Potential for Iris Darkening17.6 Potential for Unexpected HairGrowth or Eyelash Changes17.7 When to Seek Physician Advice17.8 Use with Contact Lenses17.9 FDA-Approved Patient PackageInsert*Sections or subsections omitted from the full prescribing information are not listed.FULL PRESCRIBING INFORMATION1INDICATIONS AND USAGELATISSE® (bimatoprost ophthalmic solution) 0.03% is indicated to treat hypotrichosis of the eyelashes by increasing their growth including length, thickness and darkness.2DOSAGE AND ADMINISTRATIONEnsure the face is clean, makeup and contact lenses are removed. Once nightly, place one drop of LATISSE®(bimatoprost ophthalmic solution) 0.03% on the disposable sterile applicator supplied with the package and apply evenly along the skin of the upper eyelid margin at the base of the eyelashes. The upper lid margin in the area of lash growth should feel lightly moist without runoff. Blot any excess solution runoff outside the upper eyelid margin with a tissue or other absorbent cloth. Dispose of the applicator after one use. Repeat for the opposite eyelid margin using a new sterile applicator.Do not reuse applicators and do not use any other brush/applicator to apply LATISSE®.Do not apply to the lower eyelash line (see WARNINGS AND PRECAUTIONS, 5.3 and PATIENT COUNSELING INFORMATION, 17).Additional applications of LATISSE® will not increase the growth of eyelashes.Upon discontinuation of treatment, eyelash growth is expected to return to its pre-treatment level.3DOSAGE FORMS AND STRENGTHSBimatoprost ophthalmic solution 0.3 mg/mL.4 CONTRAINDICATIONS4.1 HypersensitivityLATISSE® is contraindicated in patients with hypersensitivity to bimatoprost or any other ingredient in this product.5 WARNINGS AND PRECAUTIONS5.1 Effects on Intraocular PressureBimatoprost ophthalmic solution (LUMIGAN®) lowers intraocular pressure (IOP) when instilled directly to the eye in patients with elevated IOP. In clinical trials, in patients with or without elevated IOP, LATISSE®lowered IOP, however, the magnitude of the reduction was not cause for clinical concern.In ocular hypertension studies with LUMIGAN®, it has been shown that exposure of the eye to more than one dose of bimatoprost daily may decrease the intraocular pressure lowering effect. In patients using LUMIGAN®or other prostaglandin analogs for the treatment of elevated intraocular pressure, the concomitant use of LATISSE® may interfere with the desired reduction in IOP. Patients using prostaglandin analogs including LUMIGAN®for IOP reduction should only use LATISSE® after consulting with their physician and should be monitored for changes to their intraocular pressure (see PATIENT COUNSELING INFORMATION, 17).Pigmentation5.2 IrisIncreased iris pigmentation has occurred when the same formulation of bimatoprost ophthalmic solution (LUMIGAN®) was instilled directly onto the eye. Although iridal pigmentation was not reported in clinical studies with LATISSE®, patients should be advised about the potential for increased brown iris pigmentation which is likely to be permanent.The pigmentation change is due to increased melanin content in the melanocytes rather than to an increase in the number of melanocytes. The long term effects of increased pigmentation are not known. Iris color changes seen with administration of bimatoprost ophthalmic solution may not be noticeable for several months to years. Typically, the brown pigmentation around the pupil spreads concentrically towards the periphery of the iris and the entire iris or parts of the iris become more brownish. Neither nevi nor freckles of the iris appear to be affected by treatment. Treatment with LATISSE® solution can be continued in patients who develop noticeably increased iris pigmentation.Patients who receive treatment with LATISSE® should be informed of the possibility of increased pigmentation (see PATIENT COUNSELING INFORMATION, 17).Pigmentation5.3 LidBimatoprost has been reported to cause pigment changes (darkening) to periorbital pigmented tissues and eyelashes. The pigmentation is expected to increase as long as bimatoprost is administered, but has been reported to be reversible upon discontinuation of bimatoprost in most patients.5.4 Hair Growth Outside the Treatment AreaThere is the potential for hair growth to occur in areas where LATISSE® solution comes in repeated contact with the skin surface. It is important to apply LATISSE® only to the skin of the upper eyelid margin at the base of the eyelashes using the accompanying sterile applicators, and to carefully blot any excess LATISSE® from the eyelid margin to avoid it running onto the cheek or other skin areas (see PATIENT COUNSELING INFORMATION, 17).5.5Intraocular InflammationLATISSE® solution should be used with caution in patients with active intraocular inflammation (e.g., uveitis) because the inflammation may be exacerbated.5.6Macular EdemaMacular edema, including cystoid macular edema, has been reported during treatment with bimatoprost ophthalmic solution (LUMIGAN®) for elevated IOP. LATISSE®should be used with caution in aphakic patients, in pseudophakic patients with a torn posterior lens capsule, or in patients with known risk factors for macular edema.5.7 Contamination of LATISSE® or ApplicatorsThe LATISSE® bottle must be kept intact during use. It is important to use LATISSE®solution as instructed, by placing one drop on the single-use-per eye applicator. The bottle tip should not be allowed to contact any other surface since it could become contaminated. The accompanying sterile applicators should only be used on one eye and then discarded since reuse of applicators increases the potential for contamination and infections. There have been reports of bacterial keratitis associated with the use of multiple-dose containers of topical ophthalmic products (see PATIENT COUNSELING INFORMATION, 17).5.8 Use with Contact LensesLATISSE® contains benzalkonium chloride, which may be absorbed by soft contact lenses. Contact lenses should be removed prior to application of solution and may be reinserted 15 minutes following its administration (see PATIENT COUNSELING INFORMATION, 17).REACTIONS6 ADVERSEThe following information is based on clinical trial results from a multicenter, double-masked, randomized, vehicle-controlled, parallel study including 278 adult patients for four months of treatment.The most frequently reported adverse events were eye pruritus, conjunctival hyperemia, skin hyperpigmentation, ocular irritation, dry eye symptoms, and erythema of the eyelid. These events occurred in less than 4% of patients.Adverse reactions reported with bimatoprost ophthalmic solution (LUMIGAN®) for the reduction of intraocular pressure include, ocular dryness, visual disturbance, ocular burning, foreign body sensation, eye pain, blepharitis, cataract, superficial punctuate keratitis, eye discharge, tearing, photophobia, allergic conjunctivitis, asthenopia, increases in iris pigmentation, conjunctival edema, abnormal hair growth, iritis, infections (primarily colds and upper respiratory tract infections), headaches, and asthenia.8 USE IN SPECIFIC POPULATIONS8.1 PregnancyPregnancy Category CTeratogenic effects:In embryo/fetal developmental studies in pregnant mice and rats, abortion was observed at oral doses of bimatoprost which achieved at least 33 or 97 times, respectively, the maximum intended human exposure (based on blood AUC levels after topical ophthalmic administration to the cornea or conjunctival sac).At doses at least 41 times the maximum intended human exposure, the gestation length was reduced in the dams, the incidence of dead fetuses, late resorptions, peri- and postnatal pup mortality was increased, and pup body weights were reduced.There are no adequate and well-controlled studies of bimatoprost ophthalmic solution 0.03% administration in pregnant women. Because animal reproductive studies are not always predictive of human response, LATISSE® should be administered during pregnancy only if the potential benefit justifies the potential risk to the fetus.Mothers8.3 NursingIt is not known whether LATISSE® solution is excreted in human milk, although in animal studies, bimatoprost has been shown to be excreted in breast milk. Because many drugs are excreted in human milk, caution should be exercised when LATISSE® is administered to a nursing woman.Use8.4 PediatricSafety and effectiveness in pediatric patients have not been established.Use8.5 GeriatricNo overall clinical differences in safety or effectiveness have been observed between elderly and other adult patients.11 DESCRIPTIONLATISSE® (bimatoprost ophthalmic solution) 0.03% is a synthetic prostaglandin analog. Its chemical name is (Z)-7-[(1R,2R,3R,5S)-3,5-Dihydroxy-2-[(1E,3S)-3-hydroxy-5-phenyl-1-pentenyl]cyclopentyl]-N-ethyl-5-heptenamide, and its molecular weight is 415.58. Its molecular formula is C25H37NO4.Its chemical structure is:2H5water. LATISSE® is a clear, isotonic, colorless, sterile ophthalmic solution with an osmolality of approximately 290 mOsmol/kg.Contains:Active: bimatoprost 0.3 mg/mL; Preservative: benzalkonium chloride 0.05 mg/mL; Inactives: sodium chloride; sodium phosphate, dibasic; citric acid; and purified water. Sodium hydroxide and/or hydrochloric acid may be added to adjust pH. The pH during its shelf life ranges from 6.8 - 7.8.12 CLINICALPHARMACOLOGY12.1 Mechanism of ActionBimatoprost is a structural prostaglandin analog. Although the precise mechanism of action is unknown the growth of eyelashes is believed to occur by increasing the percent of hairs in, and the duration of the anagen or growth phase.12.3 PharmacokineticsAbsorptionAfter one drop of bimatoprost ophthalmic solution 0.03% was administered once daily into both eyes (cornea and/or conjunctival sac) of 15 healthy subjects for two weeks, blood concentrations peaked within 10 minutes after dosing and were below the lower limit of detection (0.025 ng/mL) in most subjects within 1.5 hours after dosing. Mean C max and AUC0-24hr values were similar on days 7 and 14 at approximately 0.08 ng/mL and0.09 ng•hr/mL, respectively, indicating that steady state was reached during the first week of ocular dosing. There was no significant systemic drug accumulation over time.DistributionBimatoprost is moderately distributed into body tissues with a steady-state volume of distribution of 0.67 L/kg. In human blood, bimatoprost resides mainly in the plasma. Approximately 12% of bimatoprost remains unbound in human plasma.MetabolismBimatoprost is the major circulating species in the blood once it reaches the systemic circulation. Bimatoprost then undergoes oxidation, N-deethylation, and glucuronidation to form a diverse variety of metabolites.EliminationFollowing an intravenous dose of radiolabeled bimatoprost (3.12 µg/kg) to six healthy subjects, the maximum blood concentration of unchanged drug was 12.2 ng/mL and decreased rapidly with an elimination half-life of approximately 45 minutes. The total blood clearance of bimatoprost was 1.5 L/hr/kg. Up to 67% of the administered dose was excreted in the urine while 25% of the dose was recovered in the feces.TOXICOLOGY13 NONCLINICAL13.1 Carcinogenesis, Mutagenesis, Impairment of FertilityBimatoprost was not carcinogenic in either mice or rats when administered by oral gavage at doses of up to 2 mg/kg/day and 1 mg/kg/day respectively (approximately 192 and 291 times the recommended human exposure based on blood AUC levels after topical corneal and/or conjunctival sac administration respectively) for 104 weeks.Bimatoprost was not mutagenic or clastogenic in the Ames test, in the mouse lymphoma test, or in the in vivo mouse micronucleus tests.Bimatoprost did not impair fertility in male or female rats up to doses of 0.6 mg/kg/day.14CLINICAL STUDIESLATISSE® solution was evaluated for its effect on overall eyelash prominence in a multicenter, double-masked, randomized, vehicle-controlled, parallel study including 278 adult patients for four months of treatment. The primary efficacy endpoint in this study was an increase in overall eyelash prominence as measured by at least a 1-grade increase on the 4-point Global Eyelash Assessment (GEA) scale, from baseline to the end of the treatment period (week 16). LATISSE®was more effective than vehicle as measured by theGEA score, with statistically significant differences seen at 8-week, 12-week, and 16-week (primary endpoint) treatment durations.Table 1Number (%) of subjects with at least a 1-grade increase from baseline in Global Eyelash Assessment (Primary Efficacy Endpoint – Week 16)Week LATISSE®N=137N (%) Vehicle N=141 N (%)1 7 (5%) 3 (2%)4 20 (15%) 11 (8%)8 69 (50%) 21 (15%)12 95 (69%) 28 (20%)16 107 (78%) 26 (18%)20 103 (79%) 27 (21%)In this study, patients were also evaluated for the effect of LATISSE® solution on the length, thickness and darkness of their eyelashes. Improvements from baseline in eyelash growth as measured by digital image analysis assessing eyelash length, fullness/thickness, and darkness were statistically significantly more pronounced in the bimatoprost group at weeks 8, 12, and 16.Table 2Efficacy endpoint atWeek 16(Mean Change fromBaseline)LATISSE®VehicleEyelash growth (length) (mm; % increase) N=1371.4; 25%N=1410.1; 2%Fullness/thickness (mm2; % increase) N=1360.7; 106%N=1400.1; 12%Eyelash darkness (intensity*;% increase in darkness) N=135-20.2; -18%N=138-3.6; -3%* a negative value is representative of eyelash darkeningAfter the 16-week treatment period, a 4-week post-treatment period followed during which the effects of bimatoprost started to return toward baseline. The effect on eyelash growth is expected to abate following longer term discontinuation.16 HOW SUPPLIED/STORAGE AND HANDLINGLATISSE® (bimatoprost ophthalmic solution) 0.03% is supplied sterile in opaque white low density polyethylene dispenser bottles and tips with turquoise polystyrene caps accompanied by 60 sterile, disposable applicators:3 mL in a 5 mL bottle NDC 0023-3616-03Storage: LATISSE® should be stored at 2° to 25°C (36° to 77°F).INFORMATIONCOUNSELING17 PATIENT17.1Nightly ApplicationPatients should be informed that LATISSE® (bimatoprost ophthalmic solution) should be applied every night using only the accompanying sterile applicators. They should start by ensuring their face is clean, all makeup is removed, and their contact lenses removed (if applicable). Then carefully place one drop of LATISSE®solution on the disposable sterile applicator and brush cautiously along the skin of the upper eyelid margin at the base of the eyelashes. If any LATISSE® solution gets into the eye proper, it will not cause harm. The eye should not be rinsed.Additional applications of LATISSE® will not increase the growth of eyelashes.Patients should be informed not to apply to the lower eyelash line. Any excess solution outside the upper eyelid margin should be blotted with a tissue or other absorbent material.The onset of effect is gradual but is not significant in the majority of patients until 2 months. Patients should be counseled that the effect is not permanent and can be expected to gradually return to the original level upon discontinuation of treatment with LATISSE®.17.2Handling the Bottle and ApplicatorPatients should be instructed that the LATISSE® bottle must be maintained intact and to avoid allowing the tip of the bottle or applicator to contact surrounding structures, fingers, or any other unintended surface in order to avoid contamination of the bottle or applicator by common bacteria known to cause ocular infections. Patients should also be instructed to only use the applicator supplied with the product once and then discard since reuse could result in using a contaminated applicator. Serious infections may result from using contaminated solutions or applicators.17.3 Potential for Intraocular Pressure EffectsLATISSE® may lower intraocular pressure although not to a level that will cause clinical harm.In patients using LUMIGAN® or other prostaglandin analogs for the treatment of elevated intraocular pressure, the concomitant use of LATISSE® may interfere with the desired reduction in IOP. Patients using prostaglandin analogs for IOP reduction should only use LATISSE® after consulting with their physician.17.4 Potential for Eyelid Skin DarkeningPatients should be informed about the possibility of eyelid skin darkening, which may be reversible after discontinuation of LATISSE®.17.5 Potential for Iris DarkeningAlthough iridal pigmentation was not reported in clinical studies with LATISSE®, patients should be advised about the potential for increased brown iris pigmentation which is likely to be permanent. Increased iris pigmentation has occurred when the same formulation of bimatoprost ophthalmic solution (LUMIGAN®) was instilled directly onto the eye.17.6Potential for Unexpected Hair Growth or Eyelash ChangesPatients should be informed of the possibility of hair growth occurring outside of the target treatment area if LATISSE® repeatedly touches the same area of skin outside the treatment area. They should also be informed of the possibility of disparity between eyes in length, thickness, pigmentation, number of eyelashes or vellus hairs, and/or direction of eyelash growth. Eyelash changes are likely reversible upon discontinuation of treatment.17.7When to Seek Physician AdvicePatients should be advised that if they develop a new ocular condition (e.g., trauma or infection), experience a sudden decrease in visual acuity, have ocular surgery, or develop any ocular reactions, particularly conjunctivitis and eyelid reactions, they should immediately seek their physician’s advice concerning the continued use of LATISSE®.Patients on IOP-lowering medications should not use LATISSE® without prior consultation with their physician.17.8 Use with Contact LensesPatients should be advised that LATISSE® solution contains benzalkonium chloride, which may be absorbed by soft contact lenses. Contact lenses should be removed prior to application of LATISSE® and may be reinserted 15 minutes following its administration.----------------Cut Here- --------------------------------------------------------------------------------------------------------17.9 FDA-Approved Patient Package InsertPATIENT INFORMATIONLATISSE® (la teece) (bimatoprost ophthalmic solution) 0.03%Read the Patient Information that comes with LATISSE® before you start using it and each time you get a refill. There may be new information. This leaflet does not take the place of talking with your physician about your treatment.What is hypotrichosis of the eyelashes?Hypotrichosis is another name for having inadequate or not enough eyelashes.What is LATISSE® solution?LATISSE®solution is a prescription treatment for hypotrichosis used to grow eyelashes, making them longer, thicker and darker.Who should NOT take LATISSE®?Do not use LATISSE® solution if you are allergic to one of its ingredients.Are there any special warnings associated with LATISSE® use?LATISSE®solution is intended for use on the skin of the upper eyelid margins at the base of the eyelashes. Refer to Illustration 2 below. DO NOT APPLY to the lower eyelid. If you are using LUMIGAN® or other products in the same class for elevated intraocular pressure (IOP), or if you have a history of abnormal IOP, you should only use LATISSE® under the close supervision of your physician.LATISSE® use may cause darkening of the eyelid skin which may be reversible. LATISSE® use may also cause increased brown pigmentation of the colored part of the eye which is likely to be permanent.It is possible for hair growth to occur in other areas of your skin that LATISSE® frequently touches. Any excess solution outside the upper eyelid margin should be blotted with a tissue or other absorbent material to reduce the chance of this from happening. It is also possible for a difference in eyelash length, thickness, fullness, pigmentation, number of eyelash hairs, and/or direction of eyelash growth to occur between eyes. These differences, should they occur, will usually go away if you stop using LATISSE®.Who should I tell that I am using LATISSE®?You should tell your physician you are using LATISSE® especially if you have a history of eye pressure problems.You should also tell anyone conducting an eye pressure screening that you are using LATISSE®.What should I do if I get LATISSE® in my eye?LATISSE®solution is an ophthalmic drug product. LATISSE® is not expected to cause harm if it gets into the eye proper. Do not attempt to rinse your eye in this situation.What are the possible side effects of LATISSE®?The most common side effects after using LATISSE® solution are an itching sensation in the eyes and/or eye redness. This was reported in approximately 4% of patients. LATISSE®solution may cause other less common side effects which typically occur on the skin close to where LATISSE® is applied, or in the eyes. These include skin darkening, eye irritation, dryness of the eyes, and redness of the eyelids.If you develop a new ocular condition (e.g., trauma or infection), experience a sudden decrease in visual acuity, have ocular surgery, or develop any ocular reactions, particularly conjunctivitis and eyelid reactions, you should immediately seek your physician’s advice concerning the continued use of LATISSE®solution.What happens if I stop using LATISSE®?If you stop using LATISSE®, your eyelashes are expected to return to their previous appearance over several weeks to months.Any eyelid skin darkening is expected to reverse after several weeks to months.Any darkening of the colored part of the eye known as the iris is NOT expected to reverse and is likely permanent.How do I use LATISSE®?LATISSE® solution is packaged as a 3 mL bottle of solution with 60 accompanying sterile, disposable applicators. The recommended dosage is one application nightly to the skin of the upper eyelid margin at the base of the eyelashes only.Once nightly, start by ensuring your face is clean, makeup and contact lenses are removed. Remove an applicator from its tray. Then, holding the sterile applicator horizontally, place one drop of LATISSE® on the area of the applicator closest to the tip but not on the tip (see Illustration 1). Then immediately draw the applicator carefully across the skin of the upper eyelid margin at the base of the eyelashes (where the eyelashes meet the skin) going from the inner part of your lash line to the outer part (see Illustration 2). Blot any excess solution beyond the eyelid margin. Dispose of the applicator after one use.Repeat for the opposite upper eyelid margin using a new sterile applicator. This helps minimize any potential for contamination from one eyelid to another.Illustration 1Illustration 2DO NOT APPLY in your eye or to the lower lid. ONLY use the sterile applicators supplied with LATISSE®to apply the product. If you miss a dose, don’t try to “catch up.” Just apply LATISSE® solution the next evening. Fifty percent of patients treated with LATISSE® in a clinical study saw significant improvement by 2 months after starting treatment.If any LATISSE® solution gets into the eye proper, it is not expected to cause harm. The eye should not be rinsed.Don’t allow the tip of the bottle or applicator to contact surrounding structures, fingers, or any other unintended surface in order to avoid contamination by common bacteria known to cause infections.Contact lenses should be removed prior to application of LATISSE® and may be reinserted 15 minutes following its administration.Use of LATISSE® more than once a day will not increase the growth of eyelashes more than use once a day. Store LATISSE® solution at 36o to 77o F (2o to 25o C).General Information about LATISSE®.Prescription treatments are sometimes prescribed for conditions that are not mentioned in patient information leaflets. Do not use LATISSE® solution for a condition for which it was not prescribed. Do not give LATISSE® to other people. It may not be appropriate for them to use.This leaflet summarizes the most important information about LATISSE®solution. If you would like more information, talk with your physician. You can also call Allergan’s product information department at 1-800-433-8871.What are the ingredients in LATISSE®?Active ingredient: bimatoprostInactive ingredients: benzalkonium chloride; sodium chloride; sodium phosphate, dibasic; citric acid; and purified water. Sodium hydroxide and/or hydrochloric acid may be added to adjust pH. The pH during its shelf life ranges from 6.8 - 7.8.© 2009 Allergan, Inc.Irvine, CA 92612® marks owned by Allergan, Inc.U.S. Patents 6,403,649; 7,351,404; and 7,388,029。

The Halo Beaming Model for Gamma-Ray Bursts

a rXiv:as tr o-ph/972182v121Fe b1997The Halo Beaming Model for Gamma-Ray Bursts R.C.Duncan 1&Hui Li 2ReceivedABSTRACTWe consider a model for gamma-ray bursts(GRBs)from high-velocity neutron stars in the galactic halo.In this model,bursters are born in the galactic disk with large recoil velocities V r,and GRBs are beamed to within emission cones of half-angleφb centered on V r.We describe scenarios for magnetically-channeled GRBs that have such beaming characteristics.We then make detailed comparisons of this halo beaming model(HBM)to data from the3rd BATSE Catalog and from the Pioneer Venus Orbiter experiment,for both GRB intensity and angular position distributions.Acceptableﬁts to observations of over1000bursts are obtained forφb=15◦−30◦and for a BATSE sampling depth of D∼180kpc,which corresponds to a peak burst luminosity of∼1040ergs s−1.Present data favor a truly isotropic(cosmological) model over the HBM,but not by a statistically compelling margin(<∼2σ).The HBM makes the distinctive prediction that the galactocentric quadrupole moment cos2Θ −1/3for bright,nearby GRBs is large,even though the dipole moment cosΘ remains near zero.Bursters born in nearby external galaxies,such as M31,are almost entirely undetectable in the HBM because of misdirected beaming.We analyze several reﬁnements of the basic HBM: gamma-ray intensities that vary with angle from the beam axis;non-standard-candle GRB luminosity functions;and models including a subset of bursters that do not escape from the galaxy.We also discuss the energy budgets for the bursters,the origins of their recoils,and the physics of burst beaming and alignment.One possible physical model is based on the magnetar model of soft gamma repeaters(SGRs).Empirical bounds on the rate of formation and peculiar velocities of SGRs imply that there exist∼104to∼107aged SGRs in the galactic halo within a distance of100kpc.The HBM gives an acceptableﬁt to observations only if it satisﬁes some special conditions(φb≈20◦,uniform bursting rate)which are possible,but for which there are no clear and compelling theoretical justiﬁcations.The cosmological burst hypothesis is more generic and thus more attractive in this sense.Subject headings:gamma rays:bursts—stars:magneticﬁelds—stars:neutron —galaxy:halo1.IntroductionThe Burst and Transient Source Experiment(BATSE)on the NASA Compton Gamma-Ray Observatory has revealed a nearly isotropic but inhomogeneous sky distribution of gamma-ray bursts(Meegan et al.1992;Meegan et al.1996).An additional major constraint on GRB theories,revealed by the Pioneer Venus Orbiter(PVO)experiment,is that the intensity distribution of bright bursts is consistent with a locally-uniform density of sources in Euclidean space(Fenimore et al.1993).These facts are simply and generically accounted for if the faintest observed GRBs come from cosmological distances(e.g.,Paczy´n ski1995). An alternative possibility,that we will focus on here,is that the bursts come from the extended halo or corona of our galaxy.Galactic halo models for GRBs wereﬁrst considered and tested against data by Fishman et al.(1978)and Jennings&White(1980).Reasons for favoring halo models have been summarized by Lamb(1995).In galactic halo GRB models,it is often supposed that the bursters are neutron stars, since such compact stars are conjectured to be capable of producing intenseﬂuxes of hard, non-thermal photons;while their small size could drive burst variability on submillisecond time scales.Other reasons for favoring neutron stars include observations of spectral lines, which are at present controversial,and a possible connection of classic GRBs with the March5,1979gamma ray burster(e.g.,Duncan&Thompson1992,hereafter DT92)which was localized to an angular position lying within a young supernova remnant(Cline et al.1982;Rothschild,Kulkarni&Lingenfelter1994).The displacement of this burster from the center of the supernova remnant indicates that it acquired a velocity∼1000km s−1at birth(DT92).Neutron stars with such velocities will escape the galactic disk,and move into the halo.Recent analysis of the1979March5th event lends support to its association with classical GRBs(Fenimore,Klebesadel,&Laros1996,hereafter FKL).In addition, many high-velocity radio pulsars(V>500km s−1)have been observed(Lyne&Lorimer1994;Frail1996).For these reasons,we will focus on theories of GRBs from high-velocity neutron stars(HVNSs)born in the galactic disk,asﬁrst suggested by Shklovskii&Mitrofanov (1985).In particular,we will consider the“halo beaming model”(HBM)proposed earlier (Duncan,Li,&Thompson1993,hereafter DLT;Li,Duncan,&Thompson1994,LDT;Li &Duncan,1996a,1996b;Bulik&Lamb1996).Alternative models for GRBs from HVNSs in the galactic halo invoke a delayed onset of bursting activity at a time∼107years after birth in the disk(Li&Dermer1992,hereafter LD92),which has been applied in several diﬀerent physical contexts(Colgate&Leonard1994;Lamb,Bulik&Coppi1996;Woosley &Herant1996);and weakly-bound bursters orbiting in a nonspherical galactic potential (Podsiadlowski,Rees&Ruderman1995,hereafter PRR).Models for bursters born in the galactic halo(Eichler&Silk1992;Hartmann1992;Salpeter&Wasserman1993;Wasserman &Salpeter1994),or the Magellanic Clouds(Fabian&Posiadlowski1993),are beyond the scope of this paper.In§2,we will describe our basic model for the galactic population of bursters,and brieﬂy review its physical motivations.In§3we present our basic Monte Carlo model results.We discuss how these results arise,and what they could imply for halo GRB theories.We also explore the model’s sensitivity to diﬀerent beaming angles.In§4we make detailed quantitative comparisons of the model with BATSE and PVO data.In particular, we study moments of the angular position distribution in bright subsets of observed GRBs, which is potentially a sensitive model discriminant.In§5we estimate the GRB repetition rate and the energy requirements for bursters.We also discuss two candidate power sources that could satisfy these requirements:magnetic energy and accretion energy.Severalreﬁnements of the basic HBM are investigated in§6,namely:gamma-ray intensities that vary with angular position within the beam,GRB luminosity functions,and models witha signiﬁcant subset of bursters on bound orbits in the galactic halo.In§7we give our conclusions and outline future observational tests.Note that we consider possible physical mechanisms for halo GRBs,involving neutron stars with unusually strong magneticﬁelds,in§2.3,§2.4,§5,§6.3,and Appendix B.These sections could be read separately from the rest of the paper,since they might apply in a non-beamed model context(cf.§7.2).2.The Halo Beaming Model:Physical Motivations2.1.Model AssumptionsWe will assume that GRBs are emitted by HVNSs,which emanate from the galactic disk like a“wind”extending into the galactic corona(Shklovskii&Mitrofanov1985). The HBM does not require that the bursters or their peculiar velocities have any favored orientation in a galactic coordinate frame,but it does invoke the physically-plausible condition that gamma-rays are produced only within a cone of angular radiusφb about the star’s magnetic axis±−→µ.The magnetic axis is furthermore assumed to be roughly aligned (within∼20◦)with the stellar recoil velocity V r.We discuss the physics of such beaming and alignment in§2.2.The particular version of HBM which we will analyze quantitatively in§3and§4below has the following simple properties:[1]bursters are born at positions distributed like young Pop.I stars in the galactic disk;[2]with randomly directed recoils V r=1000km s−1;[3] they emit GRBs at a constant rate,with[4]constant luminosity,and[5]the gamma ray emission is beamed parallel and anti-parallel to V r,within an angular radiusφb that we will vary.GRB beaming in a direction correlated with V r makes the observable burst distributioncomply with the BATSE dipole isotropy and the PVO brightness–distribution constraints, for the following reasons(see Figure1).Since the HVNSs are freely–streaming out of the galaxy,their mean density diminishes with distance from the galactic center as n∼r−2. However,the fraction of escaping bursters that are potentially detectable at Earth increases in proportion to the transverse area of their beaming cones,∼r2.These two trends cancell, making the eﬀective number of bursters increase linearly with the sampled volume(i.e.,an apparent“constant density”of detectable stars)within a“core radius”R c∼R o/φb,where R o∼8.5kpc is a galactic disk dimension.In the HBM,this produces the the observed “homogeneous”distribution of bright GRBs found by the PVO experiment.Furthermore, since most bursters are undetectable when they are at small r,the observable dipole anisotropy of GRB positions in the direction of the galactic center is greatly reduced.At distances larger than R c,all bursters are detectable at Earth(or nearly all;see exceptions discussed in§3),and the n∝r−2free-streaming fall-oﬀof burster density prevails, accounting for the“boundedness”( V/V max <0.5)found by BATSE.A more detailed illustration of the geometrical eﬀect of burst beaming is given in Figure2.To understand thisﬁgure,consider for a moment an idealized model in which all stars are born precisely at the galactic center(GC),and move out in random directions on straight–line trajectories,each emitting GRBs into a cone of half-opening angleφb around its velocity vector.Bursters lying outside the circles in Figure2,or within the lens-shaped intersection of circles between the GC and Earth,are then the only ones that can be detected at Earth.There is a large“zone of avoidance”(ZOA),within which all bursters are invisible.Figure2actually shows only the2-D cross-section of this ZOA.The true ZOA is the volume of revolution of the pictured shape about the line between Earth and the galactic center.The boundary in Figure2at which stars just become observable(edge of the ZOA)is part of a circle with radius D sun/sinφb,where D sun=8.5kpc is the distance from Sun to the galactic center.In the realistic situation,bursters are born throughout the galactic disk,with a peak birth rate at∼4–5kpc from the center,where the greatest concentration of Population I stars are located(van der Kruit1987).Each birthplace in the disk then has its own ZOA, scaling up or down in linear size with the distance between the birthplace and Earth.One must add the weighted distributions of detectable bursts together to get the total observable GRB distribution.There is no simple,analytic way to do this,so in what follows we will use Monte Carlo methods.We will also calculate realistic trajectories in the galactic potential, rather than assuming straight lines.2.2.Physics of Beaming and AlignmentGamma-ray bursts from high-B environments tend to be beamed alongﬁeld lines because of the transverse pair-creation opacity(e.g.,Riﬀert,M´e sz´a ros&Bagoly1989; Ho,Epstein&Fenimore1990)and because gamma rays produced by such mechanisms as curvature radiation and Compton upscattering are strongly beamed alongﬁeld lines (e.g.,Sturrock1986;Dermer1990and references therein).Even locally–isotropic emissions are strongly beamed when they occur in a relativistic outﬂow channeled by a magneticﬁeld(e.g.,Yi1993).Gamma emissions induced by sheared Alfv´e n waves in a neutron star magnetosphere can also be highly beamed(Melia&Fatuzzo1991;Fatuzzo&Melia 1993).Beaming obviates theγ–γopacity limit for GRBs(Krolik&Pier1991;Baring1993; Fenimore Epstein&Ho1993).The HBM invokes the additional(assumed)property that burster recoil velocities areapproximately aligned with magnetic dipole—and hence burst beaming—axi:−→V r −→µ.This requires that the rotation axis −→Ωis also roughly aligned with−→µ,at least to within theburst opening angleφb.Is −→V r −→µa plausible assumption?Any recoil mechanism which imparts impulses to the stellar surface with a coherence time longer than the rotation period of the star willtend to make −→V r ±−→Ωbecause of “rotational averaging.”One such recoil mechanism isanisotropic neutrino emission during the ﬁrst ∼10s after formation (§2.4below).If itis also true that −→µ ±−→Ω(at least to within the beaming angle φb )then the alignmentcondition of the HBM is satisﬁed.Such near–alignment between −→µand ±−→Ωis expected if the neutron star magnetic ﬁeldis generated by large-scale dynamo action,as in familiar stellar and planetary dynamos.The “magnetar”model,described below (§2.3),is one scenario involving such a dynamo.Vacuum magnetic torques increase alignment (Michel &Goldwire 1970;Davis &Goldstein 1970)after the star spins down past the “death line”at which magnetospheric currents are quenched,at age t D ∼105(B dipole /1014G)−1yrs (Chen &Ruderman 1993).Before this point is reached,currents might exert counter-aligning torques of comparable strength (e.g.,McKinnon 1993and references therein).Note that the GRB beaming angle φb is the sum of the intrinsic beaming angle due to gamma-ray opacity and radiative eﬀects,and the r.m.s.scatter in the angle between −→Ωand −→µ.Gamma production mechanisms such as curvature radiation and Compton upscattering could operate at signiﬁcant altitudes in the magnetosphere,where ﬁeld linesare tilted with respect to −→µ,yielding larger values of φb than one would estimate under theassumption of near-polar gamma emission (as adopted by Riﬀert,M´e sz´a ros &Bagoly 1989;Ho,Epstein &Fenimore 1990).Values as large as φb ∼20◦are possible.2.3.Magnetically-powered Bursts?The HBM might apply in a variety of diﬀerent physical contexts.Here we will brieﬂy discuss the“magnetar”model for GRBs(DT92).The magnetar idea has previously been invoked in models of soft gamma repeater bursts(Thompson&Duncan1995,hereafter TD95)as well as in models of continuous X-ray emissions from SGRs and from anomalous X-ray pulsars like1E2259+586(Thompson&Duncan1996,TD96).Magnetars are hypothetical neutron stars which are born with dipoleﬁelds in excess of B Q≡me2c3/e¯h=4.4×1013G.How might such stars form?Neutron stars are hot and convective during theﬁrst∼10seconds after they form.They also undergo strong diﬀerential rotation(TD93).Large-scale dynamo action might operate eﬃciently during this time for neutron stars with mean rotation periods below a threshold comparable to the convective overturn time,as predicted byα–Ωdynamo theory.There is evidence for such a threshold in magnetically-active main-sequence stars(e.g.,Simon1990).This threshold eﬀect could give rise to a bimodal population of neutron stars(“pulsars and magnetars”) with a factor>∼102diﬀerence inﬁeld strength(DT92).A more detailed explanation of this magnetar formation hypothesis is given by Duncan&Thompson(1996,DT96).Seven distinct estimates of theﬁeld in the bursting neutron star SGR0526−66seem to indicate B>1014G(TD95).3Magnetars spin down too rapidly to be observed as radio pulsars.DT92conjectured that GRBs areﬂare-like reconnection events in a galactic halo population of such strongly-magnetized neutron stars.We will show in§4.5that the available magnetic energy is suﬃcient to power observed GRBs,in the context of the HBM(see also PRR).Magnetic reconnection is a theoretically advantageous energy source for GRBs,for several reasons.Gravitational and nuclear energy releases generally occur in bulk baryonic matter,which has many degrees of freedom into which energy can thermalized and degraded via adiabatic expansion(e.g.,Piran&Shemi1993).Flares in a neutron star magnetosphere, on the other hand,can be“clean,”exciting only photon and pair degrees of freedom to a ﬁrst rge scale electromotive forces induced by reconnection will generally accelerate pairs,and may produce hard,non-thermal spectra via Compton upscattering, synchrotron emission,and curvature radiation(Sturrock1986).Indeed,GRBs have many qualitative similarities to stellarﬂares.Both are transient energy releases,with chaotic variability over a wide range of time scales;and both have spectrally-hard non-thermal components(e.g.,Murphy et al.1993;Ramaty&Mandzhavidze1993).The similarities are so strong that a minor fraction of bursts in the BATSE catalog might actually be intense events from nearby,non-neutron,ﬂare stars(Liang&Li1993).Stochastic avalanche models give a goodﬁt to GRB time proﬁles(Stern&Svensson1996)and to many properties of solarﬂares(Lu et al.1993).If the V r–B dipole correlation observed in some samples of radiopulsars(e.g.,Cordes 1986;Stollman&van den Heuvel1986)is applicable for all neutron stars,then the mean magnetar recoil velocity would be>∼10times larger than the mean for radiopulsars(but see Lorimer,Lyne and Anderson1995).Several mechanisms predict unusually large recoils for magnetars,suﬃcient to propel them into the galactic halo,as explained in§2.4below.Because of the dynamo origins of magnetarﬁelds,one expects that initially −→Ω,±−→V rand±−→µare roughly aligned in magnetars.4As the star spins down,magnetospheric currents might drive some degree of misalignment;however,once the star spins down past the“death line”at time∼105(B dipole/1014G)−1yrs(Chen&Ruderman1993),further spindown certainly enforces alignment(Michel&Goldwire1970;Davis&Goldstein1970), since the magnetostatic stellar distortion in magnetars is large enough to damp nutations of −→Ωabout−→µ(Goldreich1970,DLT).Incidentally,this is probably not true in old,spun-down(P rot>4s)radiopulsars.We conclude that −→V r is likely to be roughly aligned with±−→µinold magnetars,where±−→µis also the axis of beamed gamma emissions(DLT).5 This scenario for classic GRBs assumes that magnetars retain strong dipole magnetic ﬁrge-scale magnetic instabilities that reduce B dipole(Flowers&Ruderman1979) might be suppressed by toroidalﬁeld components in the stellar interior,as expected for ﬁelds generated viaα–Ωdynamos,or perhaps by other mechanisms(TD93§14.2).Note that a dipoleﬁeld anchored in the stably–stratiﬁed liquid interior of a magnetar cannot be greatly distorted by spindown-induced crustal tectonic drift(Ruderman1991)because magnetic stresses dominate tensile stresses in the crust.This might diﬀer markedly from the circumstance in ordinary radiopulsars(Ruderman1991).The interactions of vortexlines in the neutron superﬂuid with superconductingﬂux tubes are also probably irrelevant for young magnetars,because the interiorﬁeld is probably strong enough to suppress superconductivity.A young magnetar’sﬁeld evolves predominantly via ambipolar diﬀusion in the interior and Hall fracturing in the crust(TD96).As the interiorﬁeld decays, superconductivity will appear in the mantle and perhaps the core;but this probably happens only after rapid spindown has driven most superﬂuid vortex lines out of the interior.Further discussion of magnetar evolution is given by PRR.2.4.Neutrino Magnetic RecoilsA dipole anisotropy of only∼0.03in the neutrino emission from a young,hot neutron star would impart a recoil velocity∼1000km s−1to the star(Chugai1984).Here we brieﬂy review mechanisms whereby neutrino anisotropies could be magnetically induced. Purely hydrodynamic(non-magnetic)mechanisms for producing neutron star recoils have been proposed by Janka&M¨u ller(1994),Shimizu,Yamada&Sato(1994),Burrows& Hayes(1996),and in references quoted therein.A strong magneticﬁeld aﬀects neutrino emissions from the beta processes n→p+e−νand p+e−→n¯ν,as calculated by Dorofeev,Rodionov&Ternov(1985),and from neutrino scattering processes as calculated by Vilenkin(1995).Dorofeev et al.and Vilenkin furthermore estimated the neutron star recoils resulting from these processes,in the uniform-ﬁeld idealization.This is a macroscopic manifestion of parity-nonconservation in the weak interactions.6In a realistically non-uniform magneticﬁeld,the back-reaction of magnetic stresses on convective energyﬂow,along with the neutrino opacity variations outlined above,would give rise to“neutrino starspots,”analogous to sunspots,and thus produce neutron star recoils of a magnitude estimated in DT92and§13of TD93.These calculations are based on standard Weinberg-Salam weak interaction theory. If neutrinos have mass,on the other hand,a strong magneticﬁeld will shift resonantﬂavor-changing(MSW)oscillations,thereby also inducing an anisotropy in the emergent neutrinoﬂux from a nascent neutron star(Kusenko&Segr`e1996).All of these mechanisms produce recoils which vary directly—usually linearly—with the mean magneticﬁeld strength:V r∝B.Thus if any of these mechanisms dominate, one would expect much larger recoils for magnetars than for pulsars.Several non-neutrino recoil mechanisms which also might operate more eﬃciently in magnetars than in pulsars were discussed by DT92.7To estimate neutrino magnetic recoils realistically,one must know the strength,coherence length and coherence time of the magneticﬁeld during the epoch of maximum neutrino luminosity.This is not possible at present.However,if the magneticﬁeld approaches equipartition with the free energy of diﬀerential rotation and/or the convective ﬂuid mixing—which is strongly in the MHD limit—then B>∼1016G at the neutrinosphere (TD93).Such strongﬁelds could be present when most of the neutrino energy is radiated away,even if the surface dipoleﬁeld,which is frozen-in tens of seconds later by the onset of stable stratiﬁcation in the liquid interior(Goldreich&Reisenegger1992),is smaller by a factor of∼10or∼30.Recoils V r>∼103km s−1are plausible.3.Galactic Halo Model ResultsBy integrating a large number(∼106)of neutron star trajectories in a realistic galactic potential,we have derived the sky distribution of HVNSs.In the model described here(§3 and§4),all stars have|V r|=1000km s−1;thus they move in nearly straight lines and eventually escape from the galaxy(but see§6.3).We used a Monte Carlo code developed by Li&Dermer(1992),which in turn follows many of the prescriptions of Paczy´n ski(1990)and Hartmann,Epstein&Woosley(1990). In particular,we adopt van der Kruit’s(1987)model for the spatial distribution of young Pop.I stars which spawn neutron stars in the galactic disk.The trajectories of∼106 neutron stars were numerically integrated in a realistic galactic potential(Miyamoto& Nagai1975)8including a dark halo cutoﬀat radius R H=70kpc.The escape velocity fromthe center of this model potential is V esc∼600km s−1.This increases only logarithmically with R H,reaching800km s−1for R H=200kpc,thus our results are probably insensitive to R H over its plausible range(LD92).Deviations from sphericity in the dark halo were neglected(but see PRR).We also do not include the potential of M31since stars withV r∼103km s−1are negligibly perturbed by M31within the BATSE sampling depth we found of<200kpc(§4.1).In Figure3we show HBM numerical results for several angular statistics,plotted as functions of sampling depth D about the Earth.That is,theﬁgure shows cumulative angular statistics for all detectable bursters at distances from Earth that are less than or equal to the value of D on the horizontal axis.The topmost plot of Figure3shows the galactocentric dipole moment cosΘ ,whereΘis the angle between a burst and the galactic center.The second plot shows the disk-like quadrupole, sin2b −1/3,where b is galactic latitude; sin2b <1/3implies that sources are concentrated toward the disk.The third plot shows the galactocentric quadrupole, cos2Θ −1/3.The bottom plot of Figure3 shows V/V max ,a statistic which is related to the slope of the cumulative log N—log P brightness distribution,as explained below.Within each subplot of Figure3,the various lines correspond to diﬀerent beaming anglesφb as described in theﬁgure caption.We have cut oﬀtheφb=10◦and5◦curves for D≤10kpc because our Monte Carlo sampling statistics at smaller D are too poor.For example,the number of detectable stars at D≤3kpc in our Monte Carlo model is only ∼55forφb=5◦,whereas it is∼104for in the unbeamed case(φb=90◦).Increasing the sampling depth D(moving to the right in Figure3)is tantamount to including fainter and fainter bursts.Eventually the BATSE sampling depth is reached;atthis point,if the model is toﬁt observations,all the plotted statistics must simultaneously match BATSE values to within observational uncertainty.The most recent published BATSE results9are plotted in Figure3at a value D=180kpc.Before discussing these results,we must explain the signiﬁcance of V/V max .The V/V max statistic was invented for the study of the quasars(Schmidt1968),and wasﬁrst applied to GRBs by Schmidt,Higdon&Hueter(1988).For scintillation counter experiments, V/V max is the average of(C/C min)−3/2over all bursts,where C is the peak counts(in a set time interval)and C min is the threshold for detection.C min can vary with background noise and other eﬀects(e.g.,“overwrites”);however,for a source population that is uniformly distributed in static Euclidean space, V/V max is equal to0.5 regardless of how the threshold varies;and V/V max <0.5indicates that the density of bursters diminishes with distance,for standard-candle sources.Because C min is variable, the V/V max statistic is useful in the study of inhomogeneous data sets only when trying to answer the yes/no question:“Is the observed brightness distribution consistent witha uniform density of sources distributed in Euclidean space?”(e.g.,Band1992;Petrosian 1993).In Figure3,we use V/V max in a purely illustrative way.Our model values of V/V max are ideal values that would be found by an instrument with a uniform detection threshold, i.e.,no variations in the noise or the threshold settings.Since the C min values in the BATSE catalog are not highly variable(Meegan et al.1996)preliminary comparisons with BATSE,as in the bottom panel of Figure3,will not lead us astray.However,the most accurate way to compare burst observations with theory is toﬁt the observed distribution of peak photonﬂuxes to Monte Carlo models which have been realisticallyﬁltered for detection incompleteness(e.g.,Lubin&Wijers1993).This is what we do in our actual statistical comparisons of the HBM with the BATSE catalog(§4).Note that the model curves for angular statistics in Figure3implicitly assume a detector with uniform sky coverage;i.e.,we have idealized that the detector is equally capable of detecting bursts from any location on the celestial sphere.We have corrected for the imperfect sky coverage of BATSE by shifting the data points in Figure3appropriately (see footnote9).In§4we will take the opposite approach:using the raw BATSE data and ﬁltering the Monte Carlo model to take into account BATSE’s imperfect sky coverage.What can be learned from Figure3?Beaming evidently increases V/V max to nearly 0.5for nearby bursters,D∼30kpc.This can make the model satisfy PVO constraints,as we show quantitatively below.Beaming also evidently reduces cosΘ .Theφb=20◦case ﬁts BATSE observations over an appreciable range of sampling depths D>100kpc.Note that the largest deviations from isotropy in Figure3occur for bright(i.e., small-D)subsets of the observable bursts.A distinctive signature of the HBM is thatcos2Θ −1/3is positive for bright bursts.This can be understood as follows.The(bright) bursters which become visible at Earth before reaching the remote galactic halo are the oneswhich happen to be born with beaming axis−→µ(and hence,with −→V r)pointed approximatelytoward or away from Earth(to within∼φb).Because most bursters are born within the Solar circle in the galactic disk,the brightest ones tend to be seen in that direction and toward the galactic anticenter,if they have already moved past the Earth on their way out of the galaxy.Thus cos2Θ >1/3for bright bursts,while the dipole moment cosΘ remains small.。

The use of luminescent quantum dots for optical sensing

The use of luminescent quantum dots for optical sensingJose´M.Costa-Ferna´ndez,Rosario Pereiro,Alfredo Sanz-MedelSemiconductor nanocrystals,known as‘‘quantum dots’’(QDs),have demon-strated several remarkable,attractive optoelectronic characteristics espe-cially suited to analytical applications in the(bio)chemicalﬁeld.We review progress in exploiting the attractive luminescent properties of QDs in designing novel probes for chemical and biochemical optical sensing.ª2005Elsevier Ltd.All rights reserved.Keywords:(Bio)chemical sensor;Nanostructure;Photoluminescence;Quantum dot1.IntroductionQuantum dots(QDs)are nanostructuredmaterials[1],also known as zero-dimen-sional materials,semiconductor nano-crystals or nanocrystallites.These colloidalnanocrystalline semiconductors,compris-ing elements from the periodic groupsII-VI,III-V or IV-VI,are roughly sphericaland with sizes typically in the range1–12nanometer(nm)in diameter.At suchreduced sizes(close to or smaller than thedimensions of the exciton Bohr radiuswithin the corresponding bulk material),these nanoparticles behave differentlyfrom bulk solids due to quantum-conﬁnement effects[2,3].Quantumconﬁnements are responsible for theremarkable attractive optoelectronicproperties exhibited by QDs,includingtheir high emission quantum yields,size-tunable emission proﬁles and narrowspectral bands[3,4].Moreover,theirstrong size-dependent properties result in atunability emission that leads to newapplications in science and technology.The past20years have seen intenseresearch activity in the fundamental studyof the synthesis and the photophysicalproperties of QDs[5–8].Different groupshave studied II-VI semiconductor QDs,such as CdSe or CdS nanocrystals,in orderto characterize the relationship betweensize,shape and electronic properties[2,4].However,most applications so far havefocused on their use in microelectronics and opto-electrochemistry(e.g.,light-emitting diodes,solar energy conversion or quantum cascade lasers)[3,9,10]. The application of luminescent QDs as biological labels wasﬁrst reported in1998 in two breakthrough papers[11,12].Both groups simultaneously demonstrated that highly luminescent QDs can be made water-soluble and biocompatible by surface modiﬁcation and bioconjugation. They also showed the high potential of QDs as highly sensitiveﬂuorescent bio-markers and(bio)chemical probes.Other key advances enabling the emerging practical applications of QDs in biochemistry and medicine included the synthesis of high-quality colloidal QDs in large quantities[13]or recent advances on surface chemistry of QDs by conjuga-tion with appropriate functional molecules [14].The surface modiﬁcation of QDs can increase their luminescent quantum yields [14],improve stability of the nanocrystals and prevent them from aggregating[15], and make QDs available for interactions with target analytes[16],all of crucial interest for chemical sensor or biosensor applications.This article deals with work on the analytical applications of QDs in develop-ing novel(bio)chemical sensors,an area of growing interest in the past few years.We include a brief discussion on the attractive optical properties of QDs and on the importance of adequate control of the synthesis and surface modiﬁcation of the luminescent QDs,in order to achieve the desired selectivity and sensitivity for sensing target analytes.2.Optical propertiesStudies of the physical properties of QDs have revealed that strong conﬁnement ofJose´M.Costa-Ferna´ndez,Rosario Pereiro,Alfredo Sanz-Medel*Department of Physical andAnalytical Chemistry,University of Oviedo,c/Julia´nClaverı´a,8,E-33006Oviedo,Spain*Corresponding author.Tel.:+34985103474;Fax:+34985103125;E-mail:asm@uniovi.es0165-9936/$-see front matterª2005Elsevier Ltd.All rights reserved.doi:10.1016/j.trac.2005.07.008207excited electrons and holes in these nanocrystals exists at such reduced sizes and led to observations of unique optical and electronic properties[2,3].These com-pounds,which are usually non-ﬂuorescing,develop an intense,long-lasting luminescent emission when syn-thesized on an nm scale.Semiconductor QDs are char-acterized by a band-gap between their valence and conduction electron bands.When a photon having an excitation energy exceeding the semiconductor band-gap is absorbed by a QD,electrons are promoted from the valence band to the high-energy conduction band.The excited electron may then relax to its ground state by the emission of another photon with energy equal to the band-gap[4].The results of quantum conﬁnement are that the electron and hole energy states within the nanocrystals are discrete,but the electron and hole energy levels(and therefore the band-gap)is a function of the QD diameter as well as composition[17].The band-gap of semicon-ductor nanocrystals increases as their size decreases, resulting in shorter emission wavelengths[18,19].This effect is analogous to the quantum mechanical‘‘particle in a box’’,in which the energy of the particle increases as the size of the box decreases.The size-dependent emission is probably the most striking and the most studied optical property of QDs.As the emission properties of semiconductor nanocrystals depend strongly upon the energy and the density of the electron states,they can be altered by engineering the size and the shape of these tiny structures.For example,Fig.1shows theﬂuorescence spectra of CdSe QDs with different nanoparticle diameter sizes.As can be seen, differently-sized CdSe nanocrystals can be tuned in the 500–700-nm range.Moreover,as each material has tunability limits,which depend on the physical limita-tions of the dot size,other materials have been employed in QD synthesis(see Table1)(e.g.,Zn-based QDs emit below400nm while Pb-based QDs have an emission in the near-infrared spectral region).As a result of their discrete,atom-like electronic structure,QDs have typically very narrow emission spectra with full width at half-maximum(FWHM)of the luminescent emission of around15–40nm.(QDs with bandwidths as narrow as12.7–16.9nm FWHM have been reported[20]).Since the emission lines are com-paratively much narrower that those of organic dyes, detection of the QDs suffers much less from cross-talk that might result from the emission of a differentﬂuo-rophore bleeding into the detection channel of theﬂuo-rophore of interest(analyte).On the other hand QDs typically exhibit higherﬂuo-rescence quantum yields than conventional organic ﬂuorophores,allowing for greater analytical sensitivity. The quantum yield of a luminophor is a function of the relative inﬂuences of radiative recombination(producing light)and non-radiative recombination mechanisms. Non-radiative recombination,which largely occurs at the nanocrystal surface,is a faster mechanism than radiative recombination and is greatly inﬂuenced by the surface chemistry.In this context,it has for example208/locate/tracbeen demonstrated that capping the nanocrystal with a shell of an inorganic wide-band semiconductor(e.g., ZnS)reduces such non-radiative deactivations and re-sults in brighter emission[21].Chan and Nie estimated that single ZnS-capped CdSe QDs are about20times brighter that single rhodamine6G molecules[12]. There is also evidence that QDs,suitably surface-derivatized for protection,have also enhanced photolu-minescent stability as compared to typicalﬂuorescent organic dyes.Several studies have demonstrated that the photoluminescence properties of CdSe nanocrystals (including the quantum yields,peak position and FWHM)did not show any detectable change upon aging in air for several months[22].Moreover,QDs were ob-served to be100times more stable that conventional organicﬂuorophors against photobleaching[12].3.Synthesis and surface chemistryProgress on the synthesis of high-quality semiconductor nanocrystals has played and is still playing a critical role in the progress of QDs applications.Lithography-based technologies have been widely used for QDs grown onto adequate substrates[23],but have mainly been re-stricted to the preparation of optoelectronic devices. However,colloidal nanocrystals with single crystalline structure and well-controlled size and size distribution can be prepared by relatively simple nanocrystal-growth processes,starting from organometallic precursors in a mixed solvent[2,3],the latter approach being more familiar to chemists.Due to the availability of precursors and the simplicity of crystallization,CdS and CdSe have been the most well-studied colloidal QDs.Murray et al.[8]reported the synthesis of high-quality Cd-chalcogenide nanocrystals using dimethylcadmium as QD precursor in the presence of a coordinating solvent at high temperatures.The most common coordinating solvents used are trioctylphos-phine oxide,trioctylphosphine and hexadecylamine (frequently used together).Such solvents,which cap the nanocrystal and stabilize its surface,determine the particle solubility in organic media and prevent irre-versible aggregation of the nanocrystals.However,QDs capped with these hydrophobic coatings are incompati-ble with aqueous assay conditions.Consequently,in order to extend theﬁeld of application of the QDs, hydrophilic capping agents must be introduced.A landmark in the development of wet chemical routes for Cd-chalcogenide nanocrystals was the use of thiols as stabilizing agents in aqueous solution[24]. Water-soluble nanoparticles were prepared by synthe-sizing thiol-capped crystalline nanoparticles in aqueous solution by using mercapto-alcohols(e.g.,2-mercap-toethanol or1-thioglicerol)and mercapto-acids(e.g., thioglycolic acid or thiolactic acid)as stabilizers[24]. In an important paper[13],PengÕs group reported the synthesis of high-quality CdTe,CdSe and CdS nano-crystals using CdO as precursor instead of Cd(CH3)2.This latter compound is toxic,unstable,explosive and expensive,rendering QD-synthesis schemes based on its use unsuitable for large-scale synthesis(due to the need of critical experimental conditions).The quality of the QDs synthesized with this new approach[13]was found to be comparable or superior to the best previously re-ported.Moreover,the reported synthesis scheme(see Fig.2)proved to be reproducible,were based on mild and simple conditions,and had great potential to be scaled up for industrial applications.In recent years,other alternative routes for synthesis of highly mono-dispersed QDs have been investigated.For example,the use of stable non-air-sensitive precursors based on selenocarbamate derivatives of Zn or Cd[25]or on the air-stable complex Cd imino-bis(diisopropylphos-phine selenide)[26]have been proposed to synthesize monodispersed luminescent QDs of comparable quality to those prepared by more conventional methods. However,to ensure efﬁcient emission,any traps for the photogenerated electron and hole should be avoided. Possible traps in QDs are generally surface atoms that are missing at least one chemical bond.The surface atoms must be optimally constructed or reconstructed and passivated with some ligands to get rid of traps. Coating nanoparticles with a different semiconductingTable1.Nanocrystal materials and range of tunabilityQD core material a Fluorescence-emission range(nm)QD-diameter range(nm) Zinc sulﬁde(ZnS)300–410–Zinc selenide(ZnSe)370–430Cadmium sulﬁde(CdS)355–490 1.9–6.7Cadmium telluride(CdTe)620–710Lead sulﬁde(PbS)700–950 2.3–9Lead selenide(PbSe)1200–2340Lead telluride(PbTe)1800–2500a Emissionﬂuorescence of core-shell QDs is also within the range given in the second column./locate/trac209material was shown to have a profound impact on the photophysics of the nanocrystalline core[2,14,21]. Deposition of a semiconductor layer with a large band-gap(Eg)relative to the core typically results in the enhancement of the QD emission due to the suppression of radiationless recombination mediated by surface states [2,21],while the degree of charge-carrier conﬁnement does not change.Conversely,an outer layer from a semiconductor with a small Eg provides an additional area of delocalization for electron and hole[2,14,21].Of course,relaxation of the conﬁnement regime results in a red shift of the spectral features.The exciting size-dependent and surface-dependent properties of nm-sized QDs have stimulated research on surface modiﬁcation of QDs,aiming to expand their practical applications.In this context,the conjugation ofa semiconductor nanoparticle with an organic molecule[27],able to interact selectively with a target molecule or (bio)chemical species,extends the area of applications from the electronic or optical devices to the biological or chemical systems,such as the preparation of non-radioactive biological labels[11,12]or chemical opto-sensors[16].4.Optical sensing with quantum dotsMore thanﬁve years have elapsed since QDs wereﬁrst proposed as stable luminescent probes in biological labeling applications.In that pioneer work,AlivisatosÕs group[11]reported a link between biomolecules and CdS or ZnS core-shell CdSe QDs via surface coating with an additional layer of silica in order to make them bio-compatible and water soluble,and established the utility of the nanocrystals for biological staining.Simulta-neously,Chan and Nie[12]linked biomolecules to water-soluble and biocompatible QDs surface-modiﬁed with mercaptoacetic acid for ultrasensitive detection at the single-dot level.They demonstrated that conjuga-tion of the QDs with appropriate immunomolecules can be used for recognition of speciﬁc antibodies or antigens by measuring the luminescence emission of the nanoparticles.Many authors have stressed the distinct advantages of QD bioconjugates over conventional organic dyes(such as rhodamine),namely greater brightness,greater sta-bility with respect to photobleaching and narrower spectral line-widths.However,QD biological labeling has been slow to emerge into common practice,partly due to the difﬁculty in producing stable QD-biomolecule com-plexes.Developments have stressed the importance of adequate surface modiﬁcations in developing lumines-cent QDs for labeling in bioanalysis and diagnostics,as tags for protein and DNA immunoassays or as biocom-patible labels for in vivo imaging studies.Several reviews have summarized the use of luminescent QDs in such biochemical applications[27–31].Moreover,the fast development and improvements in the synthesis of QDs210/locate/trachave uncovered possibilities that analytical chemists have also started to explore in developing these nanomaterials for a new generation of optical sensors based on luminescence.4.1.Fluorescence-based transductionAs the luminescence of QDs is very sensitive to the sur-face states of the QDs,it is reasonable to expect that the chemical or physical interactions between a given chemical species and the surface of the nanoparticles would result in changes in the efﬁciency of the core electron-hole recombination [32].This has been the basis of the increase in research activity on the devel-opment of novel optical sensors based on QD probes.Following this approach,Cd-based QDs have been re-ported for optical sensing of small molecules and ions (Table 2).In a pioneering work,the addition of Cd ions to a basic aqueous solution containing unpassivated CdS nanoparticles resulted in important enhancement of the luminescence quantum yield of the nanoparticles,without detectable changes in particle sizes [7].This effect was attributed to the formation of a Cd(OH)2shell on the CdS core,which effectively eliminates the non-radiative recombination of charge carriers.A similar photoluminescence-activation effect (attrib-uted to passivation of surface trap sites that are either being ‘‘ﬁlled’’or energetically moved closer to the band edges by this simple chemical process)was also induced after adding Zn and Mn ions to colloidal solutions of CdS or ZnS QDs [32,33].This behavior provided the basis for optical sensing of such metallic cations with QDs.Besides the activation effect,QD-based optical sensing quenching strategies (based on the quenching by the analyte that affects the luminescence emission of the nanoparticle)have been proposed.Quenching mecha-nisms to explain how metal ions quench ﬂuorescence of QDs include inner ﬁlter effects,non-radiative recom-bination pathways,electron-transfer processes and ion-binding interactions.Measurement of the lumines-cence-deactivation ratio of peptide-coated CdS QDs has been proposed for the optical sensing of Cu(II)and Ag(I)[34].Similarly,the effect of three different ligands (L-cysteine,thioglycerol and polyphosphate)was evalu-ated on the luminescence deactivation of water-soluble CdS QDs with respect to several cations,including Zn and Cu ions [35].This latter work was one of the ﬁrst references to the use of luminescent QDs as selective ion probes in aqueous samples.Isarov and Chrysochoos [36]observed that the addi-tion of Cu(II)perchlorate in 2-propanol to CdS nano-particles led to the binding of copper ions onto the QD surface,accompanied by rapid reduction of Cu 2+to Cu +.It was proposed that copper ions bound onto the surface of the QDs facilitate non-radiative electron/hole (e À/h +)annihilation,thus resulting in a quenching of the luminescence from the nanoparticles.It was shown thatT a b l e 2.Q D -b a s e d f l u o r e s c e n t p r o b e s f o r c h e m i c a l d e t e r m i n a t i o n o f s m a l l m o l e c u l e s a n d i o n sQ D m a t e r i a lQ D c o a t i n gA n a l y t eM a t r i xD e t e c t i o n l i m i tM e a s u r i n g s i g n a lR e f .C d SC l y -H i s -L e u -L e u -C y sC u (I I )P h o s p h a t e b u f f e r0.5l MF l u o r e s c e n c e q u e n c h i n g[34]A g (I I )C d SP o l y p h o s p h a t e C u (I I )W a t e r0.8m M Z n (I I )F l u o r e s c e n c e q u e n c h i n g[35]L -c y s t e i n e F e (I I I )0.1m M C u (I I )T h i o g l y c e r o l Z n (I I )C d S e 2-m e r c a p t o e t h a n e s u l f o n i c a c i d C u (I I )W a t e r 3.2n M F l u o r e s c e n c e q u e n c h i n g [37]C d S e -Z n S B o v i n e s e r u m a l b u m i n C u (I I )W a t e r 10n M F l u o r e s c e n c e q u e n c h i n g [38]C d S e M e r c a p t o a c e t i c a c i d +b o v i n e s e r u m a l b u m i n A g (I )W a t e r 70n M F l u o r e s c e n c e q u e n c h i n g [39]C d T e 3-m e r c a p t o p r o p i o n i c a c i d C u (I I )W a t e r 0.19n g /m L F l u o r e s c e n c e q u e n c h i n g [40]C d T e T h i o g l y c o l i c a c i d Z n (I I ),M n (I I ),N i (I I ),C o (I I )W a t e r –F l u o r e s c e n c e q u e n c h i n g -e n h a n c e m e n t [41]C d S P o l y p h o s p h a t e I ÀM e t h a n o l –F l u o r e s c e n c e q u e n c h i n g [42]C d S e T e r t -b u t y l -n -(2-m e r c a p t o e t h y l )-c a r b a m a t e C N ÀM e t h a n o l 0.1l M F l u o r e s c e n c e q u e n c h i n g [44]C d S e 2-m e r c a p t o e t h a n e s u l f o n i c a c i d C N ÀW a t e r 1.1l M F l u o r e s c e n c e q u e n c h i n g [45]C d S L -c y s t e i n e A g +W a t e r 5.0n M F l u o r e s c e n c e e n h a n c e m e n t [46]C d S e I n c o r p o r a t e d i n p o l y m e r ﬁl m s T r i e t h y l a m i n e G a s m e d i a –F l u o r e s c e n c eq u e n c h i n g -e n h a n c e m e n t [47]B e n z y l a m i n eC d S e -Z n ST h i o g l y c o l i c +o r g a n o p h o s p h o r o u s h y d r o l a s eP a r a o x o nW a t e r 10n M F l u o r e s c e n c e q u e n c h i n g[49]/locate/trac211the quenching could be employed for chemical sensing of Cu ions in the organic solution.Water-soluble CdSe QDs with their surface modiﬁed with2-mercaptoethane sulfonic acid can be used for the sensitive and selective determination of copper(II)ions in aqueous solutions,based onﬂuorescence-quenching measurements[37].In addition,based on the photoluminescence quenching of the nanocrystals,CdSe-ZnS QDs modiﬁed with bovine serum albumin(BSA)were investigated for the determination of copper[38],CdSe QDs modiﬁed with mercaptoacetic acid and BSA were assayed for the analysis of silver[39]and CdTe nanocrystals modiﬁed with mercaptopropionic acid were proposed for the determination of Cu(II)ions[40].It was observed that the change in the absorption spectra caused by Cu(II) can be reversed by the addition of EDTA,a good com-plexing agent for Cu(II)ions.Thus,the authors proposed that the interaction between Cu(II)ions and the QD surface should be of the ion-binding type.Li et al.[41]synthesized water-soluble luminescent thiol-capped CdTe QDs and investigated the effect of divalent metal ions on their photoluminescence re-sponses.They found that zinc ions enhanced the lumi-nescence emission of the QDs.However,other metals (e.g.,calcium,magnesium,manganese,nickel and cad-mium)quenched luminescence.Apart from research on QD-basedﬂuorescent sensors for ion metals,work on other chemical species(e.g.,io-dide[42]or cyanide[43])has reported quenching the emission of CdS or CdSe QDs.A polyphosphate-stabilized CdS QD was evaluated for optical sensing of iodide[42] and found strong decay of luminescence intensity(decay times10l s),brought about by the analyte.Such quenching effects[42]were attributed to innerﬁlter effects,non-radiative recombination pathways and electron-transfer processes.The strong,reversible adsorption of negatively-charged CNÀonto the QD surface,with the consequent increased location due to compression of the electron-wave function in the QDs,was used to explain the quenching effect of cyanide[43].Following this mechanism,the synthesis of red photoluminescent CdSe QDs,with their surface modiﬁed with tert-butyl-n-(2-mercaptoethyl)-carbamate,has been proposed for the selective,sensitive determination of free cyanide in methanol after a photoactivation of the QDs[44].In a further work[45],the authors reported the syn-thesis of water-soluble luminescent CdSe QDs,surface-modiﬁed with2-mercaptoethane sulfonate,for the selective determination of free cyanide in aqueous solu-tion(see Fig.3).The addition of surfactant agents to the measurement aqueous solution was found to further greatly stabilize the QDs.In this way,theﬂuorescent signals observed allowed for high sensitivity(detection limit1.1·10À6M)and also for great selectivity of the proposed cyanide detection(over many other anionic species).As can be seen,most of the methods described so far rely on the chemical sensing of small molecules and ions with QDs via analyte-induced deactivation of photolu-minescence.However,Zhu and Chen[46]proposed a method for the determination of trace levels of silver ion based on luminescence enhancement of water-soluble CdS QDs modiﬁed with L-cysteine.The authors showed detection limits as low as5.0·10À9M.They proposed that theﬂuorescence-enhancement effect could be attributed to the formation of a complex between silver ions and the RS-groups adsorbed on the surface of the modiﬁed QDs,which resulted in the creation of radiative centers at the CdS/Ag-SR complex.The interactions between some reactive gas molecules and the surface of CdSe QDs have also been exploited in developing gas-sensing technologies[47].Nazzal et al. found that the photoluminescence of CdSe QDs incor-porated into polymer thinﬁlms is reversibly enhanced or quenched by the presence of certain gases in the212/locate/tracenvironment.After QD synthesis,photostimulation was found to be necessary to obtain a stabilized emission proﬁle and to provide reliable responses to the presence of the gases.This effect,shown in Fig.4,was also re-ported by other groups[37,44,45].Although the mechanism(s)explaining this photoactivation is(are) not clear,it is thought that a reconstruction of the sur-face atoms of the nanoparticle,or an optimization of surface-ligand passivation,could lead to the observed enhancement of the measured luminescence[47].QDs were also proposed for the design of sensing assemblies for selective detection of paraoxon[48]. Water-soluble CdSe QDs,surface functionalized with thioglycolic acid,were synthesized and incorporated to-gether with organophosphorous hydrolase(OPH)in a thinﬁlm prepared by the layer-by-layer technique.The presence of paraoxon in the sample solution was de-tected by changes in the photoluminescence emission of the QDs,attributed to an interaction of the analyte with the OPH included in the sensingﬁlm,changing its conformation.Following this approach,the synthesis of(CdSe)ZnS nanocrystals and their conjugation with organophos-phorous hydrolase(through electrostatic interaction between negatively-charged QD surfaces and the posi-tively-charged protein side chain and-NH2ending groups)has been proposed in developing a biosensor to detect paraoxon,obtaining detection limits as low as 10À8M[49].The photoluminescence intensity of the OPH/QD bioconjugate was quenched in the presence of paraoxon,matching very well with the Michaelis-Menten equation.This result indicated that the quenching was caused by the conformational change in the en-zyme,which was conﬁrmed by gas-chromatography measurements.Although such a strategy of QD-surface bioconjugation has not yet been exploited frequently for sensing,it holds great potential for further developments in optical sensing with QDs.In recent years,QDs have been used as inorganic, non-speciﬁc,DNA-binding proteins that act as lumines-cent labels for different applications(e.g.,multicolor gene mapping on the nm scale)[50].Moreover,ﬂuorescence quenching of water-soluble CdSe QDs,surface modiﬁed with mercaptoacetic acid,has also been used to develop aﬂuorescence probe for rapid,sensitive determination of DNA in a neutral medium[51].The mechanism for the binding of the nucleic acids to the QDs was investigated and it was concluded that nanoparticles bind to the helix structure of the DNA in a non-intercalative way, resulting in the observed deactivation of the lumines-cence emission of the QDs.4.2.Fluorescence(or Fo¨rster)resonanceenergy-transfer-based sensorsEnergy-transfer mechanisms have been widely used in differentﬁelds and are the basis of a new generationof Figure4.Effect of photoactivation(a)ﬂuorescence spectra from CdSe quantum dots measured after different sunlight time exposures(from ref-erence[44]),(b)Pictures of a methanolic QDs solution(1)freshly prepared and(2)after3days exposed to the sunlight./locate/trac213luminescent sensors[52].In this context,the capability of tailoring(via size)QD-photoemission properties should allow efﬁcient energy transfer with a number of con-ventional organic dyes,thus suggesting the use of the nanoparticles in sensor or chemical assay applications based on photochemically-inducedﬂuorescence(or Fo¨rster)resonance-energy-transfer(FRET)mechanisms. However,the QD emission spectrum is narrower and more symmetric than the emission from conventional organicﬂuorophores,making it much easier to distin-guish the emission of the donor from that of the accep-tor.Moreover,the high quantum yields of QDs make energy transfer very efﬁcient.Several studies have already conﬁrmed that QDs are excellent candidates for use in the design of novel FRET-based strategies.As an example,speciﬁc binding of dif-ferent proteins was observed via measurements of FRET between a CdSe-ZnS QD donor,attached to one of the proteins,and some organic acceptor dyes attached to the other protein under study.In the presence of speciﬁc interactions between both proteins,strong enhancement of the acceptor-dyeﬂuorescence was observed[53].In a more fundamental study,conjugation of BSA with luminescent CdTe nanoparticles(capped with L-cysteine)resulted in a signiﬁcant increase in the CdTe ﬂuorescent emission,attributed to an efﬁcient reso-nance-energy transfer from the tryptophan moieties of the protein units to the CdTe nanoparticles acting as acceptors[54].However,despite the demonstrated favourable properties of luminescent QDs for FRET experiments,only very few studies on the synthesis ofQD bioconjugates and their applications for QD-FRET-based optical sensors have been published so far. Luminescent CdSe QDs have been used as energy do-nors in developing a competitive FRET assay for maltose [55](see Fig.5).Semiconductor nanoparticles biocon-jugated to different maltose-binding proteins,formed using a non-covalent self-assembly scheme,act as the resonance-energy-transfer donors,while non-ﬂuorescent dyes bound to cyclodextrin serve as the energy-transfer acceptors.In the absence of maltose,cyclodextrin-dye complexes occupy the protein binding sites.Energy transfer from the QDs to the dyes quenches the QD ﬂuorescence.When maltose is present,it replaces the cyclodextrin complexes,and the QDﬂuorescence recov-ers[55].This approach has been successfully employed in developing a prototype QD-based sensor for sensing maltose in solution[56].CdSe-ZnS core-shell biocompatible QDs,stabilized by mercaptopropionic acid modiﬁed with a thiolated oligo-nucleotide,have been proposed as energy donors for lighting up the dynamics of telomerization or of DNA replication occurring on the nanoparticles,using FRET to dye units incorporated into the new synthesized telomere or DNA replica[57].After addition of telome-rase,during the progression of the telomerization,the ﬂuorescence emission from the QDs at560nm decreases with the concomitant increase of the610nm emission of the dye(using400-nm excitation).Emission observed upon telomerization is attributed by the authors to FRET from the QDs to the dye molecules incorporated into the telomeric units by telomerase.The CdSe-ZnS QDs func-tionalized with M13~DNA also enabled the detection of a viral DNA by following the DNA-replication process by FRET.Results can be applied to the fast,sensitive detection of cancer cells[57].It could be also applied to the development of chip-based DNA sensors as it func-tions like logic gates,where FRET readout occurs when hybridization and replication proceed.In one application,luminescent ZnS-capped CdSe QDs, covered with mercaptoacetic acid,have been conjugated to amine-terminated molecular beacons(MBs)at the50 end for probing DNA sequences[58].Connected to the30 end of the molecular beacon,there is a quencher mole-cule[4-(40-dimethylaminophenylazo)benzoic acid,‘‘DABCYL’’].In the absence of the target DNA sequence, MBs form a hairpin structure in which the QDs and DABCYL are in such close proximity that energy from the QDs is transferred to the quencher and noﬂuorescent signal is observed.After adding the target DNA se-quence,the MB structure opens.Since QD and quencher214/locate/trac。

liquid–liquid phase separation

Biophysical Chemistry 109(2004)105–1120301-4622/04/$-see front matter ᮊ2003Elsevier B.V .All rights reserved.doi:10.1016/j.bpc.2003.10.021Cloud-point temperature and liquid–liquid phase separation ofsupersaturated lysozyme solutionJie Lu *,Keith Carpenter ,Rui-Jiang Li ,Xiu-Juan Wang ,Chi-Bun Ching a ,a a b bInstitute of Chemical and Engineering Sciences,Ayer Rajah Crescent 28,࠻02-08,Singapore 139959,Singapore aChemical and Process Engineering Center,National University of Singapore,Singapore 117576,SingaporebReceived 31July 2003;received in revised form 8October 2003;accepted 16October 2003AbstractThe detailed understanding of the structure of biological macromolecules reveals their functions,and is thus important in the design of new medicines and for engineering molecules with improved properties for industrial applications.Although techniques used for protein crystallization have been progressing greatly,protein crystallization may still be considered an art rather than a science,and successful crystallization remains largely empirical and operator-dependent.In this work,a microcalorimetric technique has been utilized to investigate liquid–liquid phase separation through measuring cloud-point temperature T for supersaturated lysozyme solution.The effects of cloud ionic strength and glycerol on the cloud-point temperature are studied in detail.Over the entire range of salt concentrations studied,the cloud-point temperature increases monotonically with the concentration of sodium chloride.When glycerol is added as additive,the solubility of lysozyme is increased,whereas the cloud-point temperature is decreased.ᮊ2003Elsevier B.V .All rights reserved.Keywords:Biocrystallization;Microcalorimetry;Cloud-point temperature;Liquid–liquid phase separation1.IntroductionKnowledge of detailed protein structure is essen-tial for protein engineering and the design of pharmaceuticals.Production of high-quality pro-tein crystals is required for molecular structure determination by X-ray crystallography.Although considerable effort has been made in recent years,obtaining such crystals is still difficult in general,and predicting the solution conditions where pro-*Corresponding author.Tel.:q 65-6874-4218;fax:q 65-6873-4805.E-mail address:lujie@.sg (J.Lu ).teins successfully crystallize remains a significant obstacle in the advancement of structural molecu-lar biology w 1x .The parameters affecting protein crystallization are typically reagent concentration,pH,tempera-ture,additive,etc.A phase diagram can provide the method for quantifying the influence of solu-tion parameters on the production of crystals w 2,3x .To characterize protein crystallization,it is neces-sary to first obtain detailed information on protein solution phase behavior and phase diagram.Recently physics shows that there is a direct relationship between colloidal interaction energy106J.Lu et al./Biophysical Chemistry109(2004)105–112and phase diagram.Gast and Lekkerkerker w4,5x have indicated that the range of attraction between colloid particles has a significant effect on the qualitative features of phase diagram.A similar relationship should hold for biomacromolecules, i.e.the corresponding interaction potentials govern the macromolecular distribution in solution,the shape of the phase diagram and the crystallization process w6x.Many macromolecular crystallizations appear to be driven by the strength of the attractive interactions,and occur in,or close to,attractive regimes w7,8x.Recent intensive investigation has revealed that protein or colloidal solution possesses a peculiar phase diagram,i.e.liquid–liquid phase separation and sol–gel transition exists in general in addition to crystallization w9,10x.The potential responsible for the liquid–liquid phase separation is a rather short range,possibly van der Waals,attractive potential w11,12x.The measurement of cloud-point temperature T can provide useful informationcloudon the net attractive interaction between protein molecules,namely,the higher the cloud-point tem-perature,the greater the net attractive interaction. Herein Taratuta et al.w13x studied the effects of salts and pH on the cloud-point temperature of lysozyme.Broide et al.w14x subsequently meas-ured the cloud-point temperature and crystalliza-tion temperature for lysozyme as a function of salt type and concentration.From these works the cloud-point temperature was found to be typically 15–458C below the crystallization temperature. Furthermore,Muschol and Rosenberger w15x deter-mined the metastable coexistence curves for lyso-zyme through cloud-point measurements,and suggested a systematic approach to promote pro-tein crystallization.In general,an effective way to determine the strength of protein interactions is to study temperature-induced phase transitions that occur in concentrated protein solutions.Liquid–liquid phase separation can be divided into two stages w11x:(1)the local separation stage at which the separation proceeds in small regions and local equilibrium is achieved rapidly;and(2) the coarsening stage at which condensation of these small domains proceeds slowly to reduce the loss of interface free energy w16x.The coexisting liquid phases both remain supersaturated but differ widely in protein concentration.The effect of a metastable liquid–liquid phase separation on crystallization remains ambiguous w17x.Molecular dynamics simulations and analyt-ical theory predict that the phase separation will affect the kinetics and the mechanisms of protein crystal nucleation w18x.tenWolde and Frenkel w19x have demonstrated that the free energy barrier for crystal nucleation is remarkably reduced at the critical point of liquid–liquid phase separation, thus in general,after liquid–liquid phase separa-tion,crystallization occurs much more rapidly than in the initial solution,which is typically too rapid for the growth of single crystal with low defect densities w15x.The determination of the location of liquid–liquid phase separation curve is thus crucial for efficiently identifying the optimum solution conditions for growing protein crystals. Microcalorimetry has the potential to be a useful tool for determining:(1)the metastable-labile zone boundary;(2)the temperature-dependence of pro-tein solubility in a given solvent;and(3)the crystal-growth rates as a function of supersatura-tion w20x.Microcalorimeters can detect a power signal as low as a few microwatts whereas standard calorimeters detect signals in the milliwatt range. Because of this greater sensitivity,samples with small heat effects can be analyzed.In addition, microcalorimetry has the advantage of being fast, non-destructive to the protein and requiring a relatively small amount of material.The present work is concerned with the analysis of the transient heat signal from microcalorimeter to yield liquid–liquid phase separation information for lysozyme solutions at pH4.8.To further examine the role of salt and additive on interprotein interactions, cloud-point temperature T has been determinedcloudexperimentally as a function of the concentrations of salt,protein and glycerol.2.Materials and methods2.1.MaterialsSix times crystallized lysozyme was purchased from Seikagaku Kogyo,and used without further107J.Lu et al./Biophysical Chemistry 109(2004)105–112purification.All other chemicals used were of reagent grade,from Sigma Chemical Co.2.2.Preparation of solutionsSodium acetate buffer (0.1M )at pH 4.8was prepared with ultrafiltered,deionized water.Sodi-um azide,at a concentration of 0.05%(w y v ),was added to the buffer solution as an antimicrobial agent.Protein stock solution was prepared by dissolving protein powder into buffer.To remove undissolved particles,the solution was centrifuged in a Sigma centrifuge at 12000rev.y min for 5–10min,then filtered through 0.22-m m filters (Mil-lex-VV )into a clean sample vial and stored at 48C for further experiments.The concentration of protein solution was determined by measuring the absorbance at 280nm of UV spectroscopy (Shi-madzu UV-2550),with an extinction coefficient of 2.64ml y (mg cm )w 21x .Precipitant stock solution was prepared by dissolving the required amount of sodium chloride together with additive glycerol into buffer.The pH of solutions was measured by a digital pH meter (Mettler Toledo 320)and adjusted by the addition of small volumes of NaOH or HAc solution.2.3.Measurement of solubilitySolubility of lysozyme at various temperatures and precipitant y additive concentrations was meas-ured at pH 4.8in 0.1M acetate buffer.Solid–liquid equilibrium was approached through both crystallization and dissolution.Dissolving lasted 3days,while the period of crystallization was over 2weeks.The supernatant in equilibrium with a macroscopically observable solid was then filtered through 0.1-m m filters (Millex-VV ).The concen-tration of diluted supernatant was determined spec-troscopically and verified by refractive meter(Kruss)until refractive index remained unchanged ¨at equilibrium state.Solubility of each sample was measured in duplicate.2.4.Differential scanning microcalorimetry Calorimetric experiments were performed with a micro-differential scanning calorimeter with anultra sensitivity,micro-DSC III,from Setaram SA,France.The micro-DSC recorded heat flow in microwatts vs.temperature,thus can detect the heat associated with phase transition during a temperature scan.The sample made up of equal volumes of protein solution and precipitant solu-tion was filtered through 0.1-m m filters to remove dust particles further.To remove the dissolved air,the sample was placed under vacuum for 3min while stirring at 500rev.y min by a magnetic stirrer.The degassed sample was placed into the sample cell of 1.0ml,and a same concentration NaCl solution was placed into the reference cell.The solutions in the micro-DSC were then cooled at the rate of 0.28C y min.After every run,the cells were cleaned by sonicating for 10–15min in several solutions in the following order:deionized water,methanol,ethanol,acetone,1M KOH and finally copious amounts of deionized water.This protocol ensured that lysozyme was completely removed from the cells.The cells were then placed in a drying oven for several hours.The rubber gaskets were cleaned in a similar manner except acetone and 1M KOH were omitted and they were allowed to dry at low temperature.3.Results and discussionA typical micro-DSC scanning experiment is shown in Fig.1.The onset of the clouding phe-nomenon is very dramatic and easily detected.The sharp increase in the heat flow is indicative of a liquid–liquid phase separation process producing a latent heat.This is much consistent with many recent investigations of the liquid–liquid phase separation of lysozyme from solution w 22,23x .In fact,such a liquid–liquid phase separation is a phase transition with an associated latent heat of demixing.In this work,the cloud-point tempera-tures at a variety of lysozyme,NaCl and glycerol concentrations are determined by the micro-DSC at the scan rate of 128C y h.3.1.Effect of protein concentrationIn semilogarithmic Fig.2we plot the solid–liquid and liquid–liquid phase boundaries for lyso-108J.Lu et al./Biophysical Chemistry 109(2004)105–112Fig.1.Heat flow of a typical micro-DSC scan of lysozyme solution,50mg y ml,0.1M acetate buffer,pH 4.8,3%NaCl.The scan rate 128C y h is chosen referenced to the experimental results of Darcy and Wiencek w 23x .Note the large deflection in the curve at approximately 4.38C indicating a latent heat resulting from demixing (i.e.liquid–liquid phase separation )process.Fig.2.Cloud-point temperature and solubility determination for lysozyme in 0.1M acetate buffer,pH 4.8:solubility (5%NaCl )(s );T (5%NaCl,this work )(d );T (5%cloud cloud NaCl,the work of Darcy and Wiencek w 23x )(*);solubility (3%NaCl )(h );T (3%NaCl )(j ).cloud Fig.3.Cloud-point temperature determination for lysozyme as a function of the concentration of sodium chloride,50mg y ml,0.1M acetate buffer,pH 4.8.zyme in 0.1M acetate buffer,pH 4.8,for a range of protein concentrations.It is worth noting that,at 5%NaCl,our experimental data of T from cloud micro-DSC are quite consistent with those from laser light scattering and DSC by Darcy and Wiencek w 23x ,with difference averaging at approx-imately 0.88C.This figure demonstrates that liquid–liquid phase boundary is far below solid–liquid phase boundary,which implies that the liquid–liquid phase separation normally takes place in a highly metastable solution.In addition,cloud-point temperature T increases with the cloud concentration of protein.3.2.Effect of salt concentrationFig.3shows how cloud-point temperature changes as the concentration of NaCl is varied from 2.5to 7%(w y v ).The buffer is 0.1M acetate (pH 4.8);the protein concentration is fixed at 50mg y ml.Over the entire range of salt concentrations studied,the cloud-point temperature strongly depends on the ionic strength and increases monotonically with the concentration of NaCl.Crystallization is driven by the difference in chemical potential of the solute in solution and in the crystal.The driving force can be simplified as w 24xf sy Dm s kT ln C y C (1)Ž.eq109J.Lu et al./Biophysical Chemistry 109(2004)105–112Fig.4.The driving force required by liquid–liquid phase sep-aration as a function of the concentration of sodium chloride,50mg y ml lysozyme solution,0.1M acetate buffer,pH 4.8.In the same way,we plot the driving force,f ,required by liquid–liquid phase separation as a function of the concentration of sodium chloride in Fig.4.At the moderate concentration of sodium chloride,the driving force required by liquid–liquid phase separation is higher than that at low or high salt concentration.As shown in Fig.3,with NaCl concentration increasing,the cloud-point temperature increases,which is in accord with the results of Broide et al.w 14x and Grigsby et al.w 25x .It is known that protein interaction is the sum of different potentials like electrostatic,van der Waals,hydrophobic,hydration,etc.The liquid–liquid phase separation is driven by a net attraction between protein molecules,and the stronger the attraction,the higher the cloud-point temperature.Ionic strength is found to have an effect on the intermolecular forces:attractions increase with ionic strength,solubility decreases with ionic strength,resulting in the cloud-point temperature increases with ionic strength.It is worth noting that,the effect of ionic strength on cloud-point temperature depends strongly on the specific nature of the ions w 13x .Kosmotropic ions bind adjacent water molecules more strongly than water binds itself.When akosmotropic ion is introduced into water,the entro-py of the system decreases due to increased water structuring around the ion.In contrast,chaotropes bind adjacent water molecules less strongly than water binds itself.When a chaotrope is introduced into water,the entropy of the system increases because the water structuring around the ion is less than that in salt-free water.This classification is related to the size and charge of the ion.At high salt concentration ()0.3M ),the specific nature of the ions is much more important w 25x .The charges on a protein are due to discrete positively and negatively charged surface groups.In lysozyme,the average distance between thesecharges is approximately 10Aw 26x .As to the salt ˚NaCl used as precipitant,Na is weakly kosmo-q tropic and Cl is weakly chaotropic w 27x .At low y NaCl concentrations,as the concentration of NaCl increases,the repulsive electrostatic charge–charge interactions between protein molecules decrease because of screening,resulting in the increase of cloud-point temperature.While at high NaCl con-centrations,protein molecules experience an attrac-tion,in which differences can be attributed to repulsive hydration forces w 14,25x .That is,as the ionic strength increases,repulsive electrostatic or hydration forces decrease,protein molecules appear more and more attractive,leading to higher cloud-point temperature.At various salt concentra-tions,the predominant potentials reflecting the driving force for liquid–liquid phase separation are different.Fig.4shows that the driving force,f ,is parabolic with ionic strength,while Grigsby et al.w 25x have reported that f y kT is linear with ionic strength for monovalent salts.The possible reasons for that difference include,their model is based on a fixed protein concentration of 87mg y ml,which is higher than that used in our study,yet f y kT is probably dependent on protein concentration,besides the solutions at high protein and salt concentrations are far from ideal solutions.3.3.Effect of glycerolFig.5compares cloud-point temperature data for 50mg y ml lysozyme solutions in absence of glycerol and in presence of 5%glycerol,respec-110J.Lu et al./Biophysical Chemistry109(2004)105–112parison of cloud-point temperatures for lysozyme at different glycerol concentrations as a function of the con-centration of sodium chloride,50mg y ml,0.1M acetate buffer, pH4.8:0%glycerol(s);5%glycerol(j).Fig.6.Cloud-point temperatures for lysozyme at different glycerol concentrations,50mg y ml lysozyme,5%NaCl,0.1M acetate buffer,pH4.8.Fig.7.Cloud-point temperature and solubility determination for lysozyme at different concentrations of glycerol in0.1M acetate buffer,5%NaCl,pH4.8:solubility(0%glycerol)(s); T(0%glycerol)(d);solubility(5%glycerol)(h);cloudT(5%glycerol)(j).cloudtively.Fig.6shows the cloud-point temperature as a function of the concentration of glycerol.The cloud-point temperature is decreased as the addi-tion of glycerol.In semilogarithmic Fig.7we plot the solid–liquid and liquid–liquid phase boundaries at dif-ferent glycerol concentrations for lysozyme in0.1 M acetate buffer,5%NaCl,pH4.8,for a range of protein concentration.This figure demonstrates that liquid–liquid and solid–liquid phase bounda-ries in the presence of glycerol are below those in absence of glycerol,and the region for growing crystals is narrowed when glycerol is added. Glycerol has the property of stabilizing protein structure.As a result,if crystallization occurs over a long period of time,glycerol is a useful candidate to be part of the crystallization solvent and is often included for this purpose w28x.In addition,glycerol is found to have an effect on the intermolecular forces:repulsions increase with glycerol concentra-tion w29x.Our experiment results of solubility and cloud-point temperature can also confirm the finding.The increased repulsions induced by glycerol can be explained by a number of possible mecha-nisms,all of which require small changes in the protein or the solvent in its immediate vicinity.The addition of glycerol decreases the volume of protein core w30x,increases hydration and the size of hydration layer at the particle surface w31,32x. In this work,we confirm that glycerol shifts the solid–liquid and liquid–liquid phase boundaries. The effect of glycerol on the phase diagram strong-111 J.Lu et al./Biophysical Chemistry109(2004)105–112ly depends on its concentration and this canprovide opportunities for further tuning of nuclea-tion rates.4.ConclusionsGrowing evidence suggests protein crystalliza-tion can be understood in terms of an order ydisorder phase transition between weakly attractiveparticles.Control of these attractions is thus keyto growing crystals.The study of phase transitionsin concentrated protein solutions provides one witha simple means of assessing the effect of solutionconditions on the strength of protein interactions.The cloud-point temperature and solubility datapresented in this paper demonstrate that salt andglycerol have remarkable effects on phase transi-tions.The solid–liquid and liquid–liquid bounda-ries can be shifted to higher or lower temperaturesby varying ionic strength or adding additives.Ourinvestigation provides further information upon therole of glycerol used in protein crystallization.Glycerol can increase the solubility,and decreasethe cloud-point temperature,which is of benefit totuning nucleation and crystal growth.In continuingstudies,we will explore the effects of other kindsof additives like nonionic polymers on phasetransitions and nucleation rates.Much more theo-retical work will be done to fully interpret ourexperimental results.AcknowledgmentsThis work is supported by the grant from theNational Natural Science Foundation of China(No.20106010).The authors also thank Professor J.M.Wiencek(The University of Iowa)for kinddiscussion with us about the thermal phenomenaof liquid–liquid phase separation.Referencesw1x A.McPherson,Current approaches to macromolecular crystallization,Eur.J.Biochem.189(1990)1–23.w2x A.M.Kulkarni, C.F.Zukoski,Nanoparticle crystal nucleation:influence of solution conditions,Langmuir18(2002)3090–3099.w3x E.E.G.Saridakis,P.D.S.Stewart,L.F.Lloyd,et al., Phase diagram and dilution experiments in the crystal-lization of carboxypeptidase G2,Acta Cryst.D50(1994)293–297.w4x A.P.Gast, C.K.Hall,W.B.Russel,Polymer-induced phase separations in non-aqueous colloidal suspensions,J.Colloid Interf.Sci.96(1983)251–267.w5x H.N.W.Lekkerkerker,W.C.K.Poon,P.N.Pusey,et al., Phase-behavior of colloid plus polymer mixtures,Euro-phys.Lett.20(1992)559–564.w6x A.Tardieu,S.Finet,F.Bonnete,Structure of the´macromolecular solutions that generate crystals,J.Cryst.Growth232(2001)1–9.w7x D.Rosenbaum,C.F.Zukoski,Protein interactions and crystallization,J.Cryst.Growth169(1996)752–758.w8x A.George,W.W.Wilson,Predicting protein crystalli-zation from a dilute solution property,Acta Cryst.D50(1994)361–365.w9x D.Rosenbaum,P.C.Zamora, C.F.Zukoski,Phase-behavior of small attractive colloidal particles,Phys.Rev.Lett.76(1996)150–153.w10x V.J.Anderson,H.N.W.Lekkerkerker,Insights into phase transition kinetics from colloid science,Nature416(2002)811–815.w11x S.Tanaka,K.Ito,R.Hayakawa,Size and number density of precrystalline aggregates in lysozyme crys-tallization process,J.Chem.Phys.111(1999)10330–10337.w12x D.W.Liu,A.Lomakin,G.M.Thurston,et al.,Phase-separation in multicomponent aqueous-protein solutions,J.Phys.Chem.99(1995)454–461.w13x V.G.Taratuta,A.Holschbach,G.M.Thurston,et al., Liquid–liquid phase separation of aqueous lysozymesolutions:effects of pH and salt identity,J.Phys.Chem.94(1990)2140–2144.w14x M.L.Broide,T.M.Tominc,M.D.Saxowsky,Using phase transitions to investigate the effect of salts onprotein interactions,Phys.Rev.E53(1996)6325–6335. w15x M.Muschol,F.Rosenberger,Liquid–liquid phase sep-aration in supersaturated lysozyme solutions and asso-ciated precipitate formation y crystallization,J.Chem.Phys.107(1997)1953–1962.w16x C.Domb,J.H.Lebowitz,Phase Separation and Critical Phenomena,Academic,London,1983.w17x D.F.Rosenbaum,A.Kulkarni,S.Ramakrishnan,C.F.Zukoski,Protein interactions and phase behavior:sen-sitivity to the form of the pair potential,J.Chem.Phys.111(1999)9882–9890.w18x O.Galkin,P.G.Vekilov,Nucleation of protein crystals: critical nuclei,phase behavior and control pathways,J.Cryst.Growth232(2001)63–76.w19x P.R.tenWolde, D.Frenkel,Enhancement of protein crystal nucleation by critical density fluctuations,Sci-ence277(1997)1975–1978.w20x P.A.Darcy,J.M.Wiencek,Estimating lysozyme crystal-lization growth rates and solubility from isothermalmicrocalorimetry,Acta Cryst.D54(1998)1387–1394.112J.Lu et al./Biophysical Chemistry109(2004)105–112w21x A.J.Sophianopoulos,C.K.Rhodes,D.N.Holcomb,K.E.vanHolde,Physical studies of lysozyme.I.Characteri-zation,J.Biol.Chem.237(1962)1107–1112.w22x Y.Georgalis,P.Umbach, A.Zielenkiewicz,et al., Microcalorimetric and small-angle light scattering stud-ies on nucleating lysozyme solutions,J.Am.Chem.Soc.119(1997)11959–11965.w23x P.A.Darcy,J.M.Wiencek,Identifying nucleation tem-peratures for lysozyme via differential scanning calorim-etry,J.Cryst.Growth196(1999)243–249.w24x M.L.Grant,Effects of thermodynamics nonideality in protein crystal growth,J.Cryst.Growth209(2000)130–137.w25x J.J.Grigsby,H.W.Blanch,J.M.Prausnitz,Cloud-point temperatures for lysozyme in electrolyte solutions:effectof salt type,salt concentration and pH,Biophys.Chem.91(2001)231–243.w26x D.Voet,J.Voet,Biochemistry,Wiley,New Y ork,1990. w27x K.D.Collins,Charge density-dependent strength of hydration and biological structure,Biophys.J.72(1997)65–76.w28x R.Sousa,Use of glycerol and other protein structure stabilizing agents in protein crystallization,Acta Cryst.D51(1995)271–277.w29x M.Farnum, C.F.Zukoski,Effect of glycerol on the interactions and solubility of bovine pancreatic trypsininhibitor,Biophys.J.76(1999)2716–2726.w30x A.Priev,A.Almagor,S.Yedgar,B.Gavish,Glycerol decreases the volume and compressibility of proteininterior,Biochemistry35(1996)2061–2066.w31x S.N.Timasheff,T.Arakawa,Mechanism of protein precipitation and stabilization by co-solvents,J.Cryst.Growth90(1988)39–46.w32x C.S.Miner,N.N.Dalton,Glycerol,Reinhold Publishing, New Y ork,1953.。

甘南高寒草甸植物群落物种多度分布特征

中国环境科学 2021,41(3)：1405~1414 China Environmental Science 甘南高寒草甸植物群落物种多度分布特征刘旻霞*,张娅娅,李全弟,李博文,孙瑞弟,宋佳颖(西北师范大学地理与环境科学学院,甘肃兰州 730070)摘要：为分析甘南高寒草甸植物群落物种多度分布格局对海拔梯度的响应,探讨物种多样性的形成规律及维持机制,采用典型样地调查法采集数据信息,并通过RAD软件程序包对实验数据进行拟合分析.结果表明:植物群落物种多样性随海拔的升高呈现先增加后降低的“中间膨胀”模式,海拔3500m处群落物种多样性达到最大.利用5个模型对海拔梯度上全部物种多度分布进行拟合发现,甘南高寒草甸植物群落物种多度分布符合生态位模型,尤其是geo 模型,利用该模型可以很好地拟合群落物种多度分布格局,同时也说明生态位分化对高寒草甸植物群落物种多度分布起主要作用.常见种多度分布的最优拟合模型为geo模型,其与全部物种的最优拟合模型一致;而稀有种多度分布的最优拟合模型为rane模型,说明常见种是影响不同海拔梯度群落物种多度分布格局的主要原因,其对群落生产力和稳定性的维持有不可替代的作用,但稀有种因其独特的功能性状也在一定程度上影响着群落结构,二者以各自不同的方式共同维系着高寒草甸的物种多样性.关键词：高寒草甸；物种多度分布；生态位模型；常见种；稀有种；生态位分化中图分类号：X176文献标识码：A 文章编号：1000-6923(2021)03-1405-10Species abundance distribution characteristics of alpine meadow plant community in Gannan. LIU Min-xia*, ZHANG Ya-ya, LI Quan-di, LI Bo-wen, SUN Rui-di, SONG Jia-ying (College of Geography and Environmental Sciences, Northwest Normal University, Lanzhou 730070, China). China Environmental Science, 2021,41(3)：1405~1414Abstract：The response of species abundance distribution (SAD) pattern of alpine meadow plant community to altitude gradient was analyzed, and the formation rule and maintenance mechanism of species diversity were discussed. In this paper, the typical sample plot survey method was used to collect data, and the experimental data were fitted and analyzed by RAD (rank abundance distribution) software package. T he results showed that the species diversity of plant community showed a “Hump relationship” pattern with the increase of altitude, and reached the maximum at 3500m. Five models were used to fit the SAD on the altitude, it was found that the overall SAD of alpine meadow plant community was in line with the niche model, especially the geo. The model could be used to fit the distribution pattern of community species abundance, and it also showed that niche differentiation plays a major role in the SAD of alpine meadow plant community. The best fitting model of common SAD was geo, which was consistent with the best fitting model of overall species, while the best fitting model of rare SAD was rane. T he above results showed that common species were the main factors affecting the SAD pattern of communities at different altitudes, and they played an irreplaceable role in maintaining the productivity and stability of communities. However, rare species also affected the community structure to a certain extent because of their unique functional characteristics, and they maintain the stability of species diversity of alpine meadow communities in different ways.Key words：alpine meadow；species abundance；niche model；common species；rare species；niche differentiation物种多度分布格局分析对于理解植物群落物种多样性的形成和维持机制具有重要意义.多度可以反映出一个物种在植物群落中占用资源的能力,不同的植物群落具有不同的多度分布格局.物种多度分布格局是多个物种不断地相互作用、相互影响的结果,通过研究多度分布格局可揭示植物群落的组成,进而了解群落中不同物种间的关系和作用机制[1].物种多度分布和物种多样性指数都是研究群落物种多样性的重要方法,物种多度通过群落内不同物种个体数量的分布情况来反映群落的物种多样性,对于认识一个群落来说,多度分布比多样性指数更能说明问题,因而成为描述群落结构及格局的有力工具[2].物种多度分布研究始于20世纪30年代,最早由Motomura[3]提出几何级数模型用于描述水生动物种多度关系.Fisher等[4]和Preston[5]分别提出对数级数模型与对数正态模型,在对鸟类、昆虫等群落的调查研究中取得了不错的拟合效果,但此两种模型在解释群落生态学过程的形成方面出现了困难,因此,基于生态位分化理论,MacArthur[6]和Sugihara[7]先后提出了断棍模型和连续断裂模型,为对数级数模型收稿日期：2020-07-30基金项目：国家自然科学基金资助项目(31760135)* 责任作者, 教授,***************1406 中国环境科学 41卷和对数正态模型提供了生态学解释.之后,Tokeshi [8]对前人的研究做了回顾与总结,进一步提出了幂分割模型并得到了广泛的应用.以Hubbell [9-10]为代表的生态学家为解释生物多样性的分布模式在不同空间尺度上的形成机制,建立了群落中性理论模型,且在热带雨林群落物种多度的研究方面得到了很好的印证.总体而言,多度格局的拟合模型自发展以来大致出现了3种类型,即统计模型、生态位理论模型和中性理论模型;由于统计模型缺乏明确的生态学意义,难以对实际数据做出明确的解释,因而生态位理论和中性理论逐渐成为解释物种多度分布格局的两大理论模型[11-12].近年来,生态位理论模型因其明确的生态学意义而在物种多度分布模式的研究中得到了广泛的应用[13-15].青藏高原是全球海拔最高的一个独特地域单元,受长期的自然演替过程和地质活动的影响,形成了世界上最大的高山生态系统(高寒草甸和草原生态系统),其中青藏高原东部的甘南高寒草甸生物资源和物种多样性相当丰富,其对青藏高原高寒草甸生态系统有着巨大的影响.但由于全球气候变化和频繁的人类活动,甘南高寒草甸草场不断退化,生物多样性锐减,生态系统结构和功能受损[16],因此,保护和恢复甘南高寒草甸生态系统物种多样性刻不容缓.纵观目前有关青藏高原植物群落物种多度分布的研究较多,但大多都集中于不同坡向[15,17]或不同取样面积[18]上,而关于垂直带谱的研究依然比较鲜见,尤其是沿海拔梯度上植物群落物种多度分布格局的相关研究.海拔作为一个重要的环境因子会影响植物的生长和分布,虽然植物的生长并不受海拔的直接影响,但海拔的变化会使温度、光照、水分等因子发生变化从而影响植物的生长及生存,植物为了适应这种改变,会通过其生理结构、形态变化等进行响应[19].基于上述背景,选取青藏高原东北部的甘南高寒草甸区植物群落作为研究对象,采用RAD 软件程序对群落物种多度分布进行模型拟合,以期回答以下3个科学问题:(1)物种多样性随海拔梯度是如何变化的?(2)群落全部物种多度分布格局随海拔的上升是如何变化的?是否可以利用生态位模型进行拟合?(3)常见种和稀有种的多度分布格局有何不同?对以上科学问题的回答可以填补关于垂直带谱上高寒草甸植物群落物种多度分布格局研究的空缺,同时也有助于我们以这样的方法探索高寒草甸植物群落物种多样性的分布规律及形成机制,揭示该地区植物群落构建过程中的资源分配模式,从而为青藏高原高寒草甸的植被恢复和生物多样性的保护提供技术支撑和理论依据. 1 材料与方法 1.1 研究区概况图1 研究区位置示意 Fig.1 Study area location map研究区地处甘肃省西南部的甘南高寒草甸区(33°06′~36º10´N,100º46′~104º44′E)(图1),是典型的高寒草甸生态系统.该地区处于青藏高原与黄土高原的过渡带,年均温2.3℃,气温南高北低,年日照时数2000~2500h 左右,年平均降雨量约600~800mm,多集中在6~9月,降雨变率大,年蒸发量约1238mm.甘南州草地总面积2.72106hm 2,占甘南总面积的70.28%,草地类型多样,其中高寒草甸类面积最大,植被以耐寒的多年生草本植物为主,主要有金露梅(Potentilla fruticosa )、矮嵩草(Kobresia humilis )、嵩3期刘旻霞等：甘南高寒草甸植物群落物种多度分布特征 1407草(Kobresia Willd )、火绒草(Leontopodium nanum )、垂穗披碱草(Elymusnu t ans Griseb )、条叶银莲花(Anemone t rullifolia )、乳白香青(Anaphalislactea maxim )、珠芽蓼(Polygonum vivipurum )、圆穗蓼(Polygonum macrophyllum )、麻花艽(Gentiana straminea )、蛇含委陵菜(Potentilla kleiniana )、羊茅(Festuca ovina )、米口袋(Gueldenstaedtia multiflora )、莓叶委陵菜(Po t en t illa fragarioides )、甘肃棘豆(Oxy t ropis kansuensis )、马先蒿(Pedicularis resupinata )、华丽龙胆(Gentiana sino -ornata )等.土壤表层有致密紧实的草皮层,腐殖质和有机质含量较高,属于典型的高寒草甸土. 1.2 样地设置2019年7月中旬在甘南高寒草甸进行野外群落学调查并采集数据.在研究区试验点内沿海拔3000m 、3250m 、3500m 、3750m 、4000m 各选3座山体,于坡度接近、坡向相同的山坡分别设置研究样地(每个海拔梯度共设置3个样地作为重复),在每个样地随机布设8个50cm×50cm 的样方,调查并记录每个样方中物种的密度、盖度、高度等基本特征并采集生物量.采用梅花五点法在每个样方内取0~20cm 深处的土样,取得土样,带回实验室进行土壤养分的测定,同时测量该海拔处的土壤温度(ST)、大气温度、坡度(土温、大气温度从8:00~ 18:00每隔1h 进行测量,每组5个重复).样地基本信息见表1.表1 样地基本信息Table 1 Basic information of sample plots海拔(m)经度(E)纬度(N)坡向坡度(°) 放牧强度是否围栏3000 101°32′28″ 33°39′44″ 北 18. 3 ± 1.6a 相同否 3250 101°57′34″ 33°47′39″ 北 18. 6 ± 2.3a 相同否 3500 101°46′26″ 33°58′50″ 北 19.2 ± 2.8b 相同否 3750 102°06′38″ 34°27′45″ 北 18.8 ± 1.7a 相同否 4000 102°52′55″ 34°5′31″ 北19.9 ± 1.7b相同否注:不同字母表示差异显著(平均值±标准误).1.3 研究方法1.3.1 全部物种多度分布模型(1)生态位重叠模型(overlapping niche model,简称over),该模型假设群落总生态位为一条棒,每个物种的多度为棒上随机两点之间的距离,且各物种多度间不存在联系,物种都是按照需求获取生存资源,这样物种间就会出现生态位重叠[20].计算公式如下:121[1]12i i i i P P +⎡⎤=−×−⎢⎥+⎣⎦ (1) (2)几何级数模型(geometric series,简称geo),该模型定义群落中第一个优势种优先占领一部分资源,第二个优势种占领剩余资源的部分,这样依次类推,直到最后一个物种没有剩余资源可占领,体现出了生态位优势占领假说.计算公式如下:()1[]1i i E P P P −=− (2) 式中:P i 为第i 个物种的多度,P 为各物种所占资源占总资源的比例,E 为群落资源总量[20].(3)断棍模型(broken stick model,简称bro),该模型认为所有物种会在群落中同时出现,且其竞争能力相近,其多度反映了竞争物种间资源的随机分配是沿着一个一维梯度进行的.其第i 个物种的多度比例期望值为:11s i x i P s x==∑ (i =1,2,3,4,…,s) (3) 式中:i 为观察群落的物种数,S 为物种总数[20].(4)随机分配模型(random assignment model,简称rane),该模型假定群落中各物种的多度之间没有联系,不存在种间竞争,一般情况下物种不能全部占用其生态位.其第i 个种的期望多度为[20]:0.510.5i s iP =− (4) (5)对数正态分布模型(log -normal distribution model,简称norm),该模型认为群落中总个体数的对数符合正态分布,则第i 个物种的多度为:()()log log i N E μσθ⎡⎤+⎣⎦= (5) 式中:N 表示物种总数,μ和σ分别表示正态分布的均值和方差,θ表示正态偏差[20].1.3.2 常见种、稀有种多度分布模型(1)几何级数模型(geometric series,简称geo),公式同上.1408 中国环境科学 41卷(2)随机分配模型(random assignment model,简称rane),公式同上.1.3.3 数据处理及模型拟合优势度检验使用RAD 软件进行物种多度拟合,采用Excel 2010进行数据处理,用Origin 2017进行绘图.以r (用模型的n 个副本的平均密度进行计算,采用最小二乘法中加入密度最大差异和物种数量的校正因子,从而增强了模型拟合优度的区分能力)、Oc (经过Ulrich 校正过的Preston 倍程分组后的χ2检验法)、CL (置信区间)值作为优势度检验的3个指标(r 、Oc 、CL 值都由RAD 拟合输出),r 和Oc 值大于0小于100,当其越接近10或等于10时拟合效果更好;置信区间CL 值一般默认为是大于或等于95%,等于100%时拟合效果最好.1.3.4 物种重要值计算物种重要值(Importance value),将其结果作为评价各物种的群落优势度、划分常见种(Z ≥0.01)与稀有种(Z <0.01)的指标[21].其计算如下:()4Z =+++相对多度相对盖度相对频度相对高度式中:Z 为每个样地每种植物的重要值. 1.3.5 物种多样性指数计算(1)Margalef 丰富度指数[22]:()()1ln D S N =− (6)(2)Simpson 指数[23]:21si i N D N =⎛⎞=⎜⎟⎝⎠∑ (7)(3)Shannon -Weiner 指数[24]:()1ln si i i H P P ==−∑ (8)式中:S 为物种数;N 为个体总数;N i 为第i 种物种个体数,i =1,2,3… S;P i 是第i 种比例多度,给定为:P i = N i /N .2 结果与分析2.1 群落物种多样性随海拔梯度的变化越复杂的群落结构,其多样性指数、丰富度指数等就越高,这不仅在水平方向上有所体现,且在垂直方向上也有较大差异.从图2可以看出,随海拔的升高,物种丰富度、Margalef 丰富度指数与Simpson 指数的变化幅度较大且变化趋势基本一致,均随海拔的升高先增大后减小;Shannon -Weiner 指数海拔3500m 最大,其次是海拔3000m,其余3个海拔的变化幅度基本相同.由此表明,群落物种多样性沿海拔梯度出现了较大的差异,中间海拔梯度(3500m)物种多样性最大.3000325035003750 40004.55.05.56.06.57.07.58.08.59.09.5海拔(m)海拔(m)M a r g a l e f 丰富度指数3000 3250 3500 375040000.70.8 0.9 1.0 1.1S i m p s o n 指数3000325035003750 40001.891.921.951.982.01S h a n n o n -W e i n e r 指数3000 3250 3500 37504000物种丰富度图2 物种多样性随海拔梯度的变化Fig.2 Variation of species diversity with altitude gradients2.2 不同海拔全部物种多度分布拟合曲线的变化全部物种多度分布曲线斜率表明(图3),随海拔的升高,曲线斜率呈先减小后增大的趋势,海拔3000m 处曲线斜率为0.059,3500m 处为0.041,4000m3期刘旻霞等：甘南高寒草甸植物群落物种多度分布特征 1409处为0.112.由此发现在海拔3500m 曲线斜率最小,曲线最为平缓,说明此群落物种分布较为均匀,各物种相对多度差异较小,且此处全部物种相对多度之和在5个海拔中最大、物种丰富度最高、多样性最大.而海拔4000m 恰恰相反.0 20 4060 8010-310-210-1100相对多度物种序列图3 高寒草甸各海拔物种多度分布曲线Fig.3 SAD curve of alpine meadow at different altitudes2.3 全部物种多度分布的模型拟合优势度检验采用over 、geo 、bro 、rane 和norm 模型对不同海拔群落的全部物种多度分布进行拟合优势度检验,结果见表2.海拔3000m 、3250m 、3750m 、4000m 全部物种多度分布均符合geo 、bro 、rane 及over 4种模型,其中geo 模型的拟合效果优于其他3种模型;norm 模型在此4个海拔处都无法拟合,其r 值、Oc 值、CL 值与检验标准差异较大.海拔3500m 处,物种多度分布符合全部5种模型,其中拟合效果最好的是geo 模型(r =10.28,Oc =9.10,CL =1);其次是bro 模型(r =14.49,Oc =10.38,CL =1),其他3个模型的拟合效果一般.表2 高寒草甸不同海拔物种多度分布模型拟合优势度检验 Table 2 Dominance test of SAD model fitting at differentaltitudes in alpine meadow海拔(m)模型rOc Clgeo 21.18 13.71 1.00 bro 41.33 15.90 1.00 3000 over 29.71 12.18 1.00 norm 35.89 43.35 0 rane 32.07 10.92 0.95 geo 12.36 16.01 1.00 bro 28.56 19.90 1.00 3250 over 24.71 12.31 1.00 norm 105.63 37.29 0 rane 54.38 32.56 0.95geo 10.28 9.10 1.00bro 14.49 10.38 1.00 3500 over 48.09 14.33 0.97 norm 28.69 5.44 1.00 rane 18.39 34.84 0.96 geo 30.18 25.37 1.00 bro 35.32 28.72 0.98 3750 over 54.68 28.42 0.96 norm 85.31 112.73 0.69 rane 47.53 29.83 1.00 geo 22.21 32.82 1.00 bro 70.42 18.20 0.96 4000 over 84.33 32.89 0.96 norm 78.18 107.33 0.73 rane 44.68 20.06 0.97 注:r 为采用最小二乘法中加入了密度最大差异和物种数量的校正因子;Oc 为Preston 倍程分组后的χ2检验法;CL 为置信区间.加黑数据为每个海拔处的最优模型拟合结果.2.4 全部物种多度拟合值和实测值的比较0 10 2030 4001020304010-310-210-110001020 30 40 50 3000m相对多度3250m相对多度3500m相对多度物种序列物种序列物种序列10-3 10-210-1 100 10-310-210-11000 1020 30 40010203010-310-210-11003750m相对多度实测值拟合值geo bro over norm rane物种序列物种序列4000m相对多度10-310-210-1 100图4 甘南高寒草甸不同海拔实测值和拟合值的对比Fig.4 Comparison of measured and fitted values of alpine meadow at different altitudes1410 中国环境科学 41卷图4表明,全部物种多度拟合值和实测值的变化趋势相似.海拔3000m、3250m、3750m和4000m 处geo、bro、over和rane 4个模型的拟合值和实测值的变化趋势基本一致,均随物种序列的增加逐渐减小;norm模型对此4个海拔群落物种多度的拟合结果与实测值的变化趋势有很大差异,实测值随物种序列的增加逐渐减小,而拟合值随物种序列的增加基本无变化,说明其拟合结果对物种多度分布只有统计学意义,而在资源分配上并无实际意义.海拔3500m处,5种模型的拟合值均随物种序列的增加逐渐减小,和实测值的变化趋势一致,说明5种模型对于该海拔群落物种的资源分配模式都具有一定的指导意义.而在5个模型中,geo模型对各海拔物种多度的拟合效果最好,其均表现出随着物种序列的增加,物种多度的拟合值逐渐小于实测值的特征,说明在资源有限的情况下,不能使每个物种都维持较高的多度,这种逐渐减少的特征也符合geo模型的分布模式.2.5 常见种和稀有种多度的模型拟合优势度检验常见种的模型拟合优势度检验结果表明,geo 模型对各个海拔常见种的拟合效果整体较好,其r 值和Oc值均小于100接近于10,置信区间CL值均接近1;但相比之下,geo模型对海拔3500m常见种多度的拟合效果明显优于其他海拔(r=10.43, Oc=10.59, CL=1).rane模型只对海拔3500m处常见种的多度分布可以拟合,但拟合效果一般(r= 20.44,Oc=59.83, CL=0.96),而对其他4个海拔无法拟合(表3).稀有种的模型拟合优势度检验结果表明,各个海拔稀有种的多度分布均符合rane模型,其中该模型对3500m的拟合效果最好(r=12.32,Oc=8.70, CL=1).geo模型只能拟合海拔3250m和3500m处稀有种的多度分布,但拟合效果都一般,而在其余海拔处无法拟合.表3 常见种和稀有种的模型拟合优势度检验Table 3 Model fitting dominance test of common and rare species物种常见种稀有种海拔(m) 模型r Oc CL r Oc CL 3000 geo 24.58 8.39 1 55.41 4.70 0 rane87.3256.17 0 36.783.13 13250 geo 31.26 18.71 1 47.25 15.72 0.98 rane65.8294.22 0 43.286.71 0.963500 geo 10.43 10.59 1 43.59 4.20 1 rane 20.44 59.83 0.96 12.32 8.70 1 3750 geo 42.58 27.13 0.98 89.05 101.23 0.87 rane116.3254.750.87 8.8337.26 14000 geo 23.46 36.59 1 33.09 115.06 0.79 rane65.28104.820.8532.494.83 1 注:r为采用最小二乘法中加入了密度最大差异和物种数量的校正因子;Oc为Preston倍程分组后的x2检验法;CL为置信区间.2.6 常见种和稀有种多度拟合值和实测值的比较就常见种来说,geo模型对其多度的拟合结果与实测值的变化趋势基本吻合,均随物种序列的增加逐渐减小(图5(a));rane模型只对海拔3500m 群落常见种多度分布有较好的拟合效果,其他4个海拔常见种多度实测值随物种序列的增加逐渐减小,而拟合值随物种序列的增加无任何变化(图5(c)),其拟合结果对常见种多度分布没有任何实际意义.由此说明常见种在环境中资源的获取方式偏向固定分配.就稀有种来说,geo模型只在海拔3250m和3500m处物种多度的拟合值和实测值的变化趋势一致,在其余海拔处拟合值与实测值差异较大,其实测值均随物种序列的增加逐渐减小,而拟合值基本无变化或变化甚微(图5(b));rane模型对各个海拔稀有种的拟合结果与实测值的变化趋势基本一致(图5(d)),其拟合效果明显优于geo模型,说明稀有种在各个海拔的资源分配模式偏向随机分配.3期刘旻霞等：甘南高寒草甸植物群落物种多度分布特征 14110 1020 30 40 50607010-310-210-11000510152025 30 35 10-210-1100(a)常见种geo 3500m4000m 3750m 3250m 3000m 相对多度(b)稀有种geo4000m 3750m 3500m 3250m 3000m相对多度实测值geo510152025 30 3510-210-11000 1020 30 40 50 60708010-310-210-1100(c)常见种rane4000m3750m (d)稀有种rane4000m 3750m 3500m 3250m 3000m 相对多度物种序列3500m 3250m 3000m 相对多度物种序列实测值rane图5 不同海拔常见种、稀有种实测值和拟合值比较Fig.5 Comparison of measured and fitted values of common and rare species at different altitudes3 讨论海拔作为重要的地形因子在很大程度上会影响植物的生长发育,不同海拔常常因其存在异质性变化的水热环境因子而引起植物性状发生适应性变化,形成具有不同生活策略的植物群落[25],因此从海拔梯度方面探讨其对群落物种多度分布的影响对于研究高寒草甸植物群落的物种多样性变化具有重要意义.3.1 物种多样性随海拔梯度的变化物种多度分布曲线的斜率变化可以判断群落的物种丰富度,进而判断群落物种多样性的变化,多度分布曲线斜率越小,则群落物种多样性越大;反之,物种多样性越小[26].本文通过研究海拔梯度上植物群落物种多样性(图2)及物种多度分布(图3)发现,甘南高寒草甸植物群落物种多样性随海拔的变化出现“中间膨胀”模式,这与其他研究结果一致[27],即物种多样性分布存在随海拔的升高先增大后减小的规律,最大值往往出现在中间海拔.其可能的解释是低海拔梯度植物群落在发育过程中,充分的环境资源使其维持着较高的物种多样性,但由于低海拔梯度受外界干扰较大,尤其是人类活动的影响会导致某些物种丧失[19],对植被的空间分布产生较大的影响,这在一定程度上会降低群落物种多样性.在中间海拔3500m 处,外界干扰变小,水热条件及土壤养分达到最优,丰富的环境资源足以支撑更多不同物种的生长发育,其物种多样性也会随之增大.当海拔超过3500m 向更高海拔过渡时,温度逐渐降低,生境越来越严酷,环境的筛选作用降低了群落物种多样性.3.2 全部物种多度分布对海拔梯度的响应由于自然群落结构的复杂性,不同物种多度分布模型表示的生态学意义更能够真实有效地反映群落的内在特征[20].甘南高寒草甸植物群落全部物种多度分布随海拔升高发生了明显的变化,导致与其符合的多度分布模型也随之改变,同时也改变了其群落构建的生态学过程(表2,图4).海拔3000m 和3250m 处,植物群落物种多度分布符合geo 、bro 、over 和rane 4种模型,其中,geo 模型的拟合效果较好,其次是over 模型,这可能是因为低海拔地区环境条件相对优越,植物赖以生存的资源较为充足,物种间的竞争作用较小,确定性的生态位划分在群落构建过程中占据主要地位;但同一群落各物种相似的生态适应性会导致产生较高的生态位重叠,尤其对于物种数量多且差异较小、物种分布比较均匀的群落来说更为显著[28-29].海拔3500m 处群落物种多度分布符合全部5种模型,其中geo 模型的拟合效果甚优,其次是bro 模型,而geo 模型和bro 模型都代表了资1412 中国环境科学 41卷源分配的确定性特征.这可能是因为此处属中海拔地段,水热条件适合更多的物种生存,因此物种丰富度较高,根据生态位补偿效应可知不同物种能更有效地获得较为充足的生存资源,相比竞争,资源共享在此处可能更为突出,所以其资源分配很大程度上属于固定分配模式,这与其它研究结果一致[30].海拔3750m处,除norm外其他4种模型均可拟合群落物种多度分布,但其拟合优势度都较低,这可能是因为该海拔处于3500m和4000m的过渡地带,逐渐减少的环境资源在一定程度上无法满足群落所有物种的发展需求,植物的生长处于不稳定的状态,导致群落物种分布格局受制于多个生态学过程.在海拔4000m处,同样是geo模型的拟合效果相对较好,可能的解释是高海拔地区环境条件恶劣,其生境因子对物种的生长要求严苛,且在此处物种赖以生存的资源有限,物种之间存在较为激烈的竞争,受这2大因素的影响,一般的物种难以在此处存活,尤其是敏感度高的稀有种[31],在长期的资源竞争中,一些弱势群体被淘汰,未被淘汰的物种各自占据着所需生态位稳定共存.由此说明,生态位模型可以很好地拟合青藏高原高寒草甸植物群落物种多度分布,尤其是geo模型,但这并不说明其群落多度分布格局一定符合geo模型所代表的生态学过程,由于geo模型划分资源的方式以确定性模式为主,因此至少可以推测出高寒草甸群落物种多度分布主要受生态位分化作用的影响,这与其他研究结果一致[15,18].3.3 常见种和稀有种多度分布随海拔梯度的变化通过常见种和稀有种的多度分布模式来探讨植物群落物种多度分布的变化是研究植物群落物种多度分布格局的一种重要且有效的方法[32].常见种多度分布模式通常趋向于群落整体多度分布模式或与之持有较高的一致性,而稀有种因其独特的功能性状对于群落物种多样性的增加及生态系统功能的维持也有突出的贡献.王世雄等[33]通过研究黄土高原辽东栎群落稀有种和常见种对物种多样性的贡献发现,稀有种和常见种对物种多样性的贡献存在空间尺度差异性,并且这种差异不依赖于稀有种和常见种的相对比例.刘旻霞等[34]研究了甘南亚高寒草甸稀有种对物种多样性和物种多度分布格局的贡献,结果表明稀有种在甘南亚高寒草甸物种多样性中的相对贡献高于非稀有种.本研究发现,常见种的最优拟合模型为geo模型,其与全部物种多度分布的最优模型一致(表3,图5a,图5c),说明其多度分布模式是引起群落物种多度分布形成的主要原因,同时也说明常见种在生境中获取资源的方式可能以固定分配为主.这是因为作为常见种的矮嵩草、嵩草等植物,它们也是高寒草甸的优势种或亚优势种,这些物种具有较强的适应严酷环境的功能性状,生存能力较强,它们对群落类型的形成和物种丰富度的主导作用受自身比例及稀有种比例等因素的影响甚微[35].就稀有种来说,rane模型可以拟合所有海拔梯度上稀有种的多度分布;geo模型除了在海拔3250m和3500m处可以拟合外,其他海拔皆无法拟合(表3,图5b,图5d).其中,rane模型的拟合效果优于geo模型,说明稀有种的资源分配模式更倾向于随机分配.这可能是因为稀有种本身对各类因素的干扰更为敏感,竞争力无法与其他常见种或者优势物种相抗衡,在激烈的竞争过程中会逐渐减少甚至丧失,其生长状态缺乏稳定性,所以对于稀有种来说,其资源获取模式可能更偏向随机分配,这与相关研究[15][36][37]结果一致.由此表明,群落内物种分布不只遵循一种规律,利用不同模型拟合群落内常见种和稀有种多度分布效果可能会更好[8].4结论4.1甘南高寒草甸植物群落物种多样性随海拔的不同出现“中间膨胀”模式.4.2物种多度分布随海拔的不同出现了较为显著的变化,海拔3500m处群落多度拟合符合geo、bro、over、rane及norm 5种模型;其余4个海拔除norm 模型外,其他4种模型均可拟合;在5个海拔梯度上,geo模型的拟合效果最佳;生态位分化在高寒草甸植物群落的构建过程中起主导作用,利用生态位模型(尤其是geo模型)可以很好地拟合高寒草甸植物群落物种多度分布.4.3常见种多度分布的最优拟合模型为geo模型,与全部物种多度分布的最优拟合模型一致,稀有种多度分布的最优拟合模型为rane模型;常见种对不同海拔梯度群落物种多度分布的影响高于稀有种,其对群落生产力和稳定性的维持有不可替代的作用,但稀有种因其独特的功能性状也在一定程度上影响着群落结构,二者以各自不同的方式共同维系。

UBVI CCD photometry of two old open clusters NGC 1798 and NGC 2192

a r X i v :a s t r o -p h /9901140v 1 12 J a n 1999Mon.Not.R.Astron.Soc.000,1–8(1998)Printed 1February 2008(MN L A T E X style ﬁle v1.4)UBV I CCD photometry of two old open clusters NGC 1798and NGC 2192Hong Soo Park and Myung Gyoon Lee ⋆Department of Astronomy,Seoul National University,Seoul 151-742,KoreaAccepted 1998??.Received 1998??;in original form 1998July ??ABSTRACT We present UBV I CCD photometry of two open clusters NGC 1798and NGC 2192which were little studied before.Color-magnitude diagrams of these clusters show several features typical for old open clusters:a well-deﬁned main-sequence,a red giant clump,and a small number of red giants.The main sequence of NGC 1798shows a distinct gap at V ≈16.2mag.From the surface number density distribution we have measured the size of the clusters,obtaining 8′.3(=10.2pc)for NGC 1798and 7′.3(=7.5pc)for NGC 2192.Then we have determined the reddening,metallicity,and distance of these clusters using the color-color diagrams and color-magnitude diagrams:E (B −V )=0.51±0.04,[Fe/H]=−0.47±0.15dex and (m −M )0=13.1±0.2(d =4.2±0.3kpc)for NGC 1798,and E (B −V )=0.20±0.03,[Fe/H]=−0.31±0.15dex and (m −M )0=12.7±0.2(d =3.5±0.3kpc)for NGC 2192.The ages of these clusters have been estimated using the morphological age indicators and the isochrone ﬁtting with the Padova isochrones:1.4±0.2Gyrs for NGC 1798and 1.1±0.1Gyrs for NGC 2192.The luminosity functions of the main sequence stars in these clusters are found to be similar to other old open clusters.The metallicity and distance of these clusters are consistent with the relation between the metallicity and galactocentric distance of other old open clusters.Key words:Hertzsprung–Russell (HR)diagram –open clusters and associations:general –open clusters and associations:individual:NGC 1798,NGC 2192–stars:luminosity function.1INTRODUCTIONOld open clusters provide us with an important informationfor understanding the early evolution of the Galactic disk.There are about 70known old open clusters with age >1Gyrs (Friel 1995).These clusters are faint in general so thatthere were few studies about these clusters until recently.With the advent of CCD camera in astronomy,the numberof studies on these clusters has been increasing.However,there are still a signiﬁcant number of old open clusters forwhich basic parameters are not well known.For example,metallicity is not yet known for about 30clusters amongthem.Recently Phelps et al.(1994)and Janes &Phelps (1994)presented an extensive CCD photometric survey of potentialold open clusters,the results of which were used in the studyon the development of the Galactic disk by Janes &Phelps(1994).In the sample of the clusters studied by Phelps et al.⋆corresponding author,E-mail:mglee@astrog.snu.ac.krthere are several clusters for which only the non-calibrated photometry is available.We have chosen two clusters among them,NGC 1798and NGC 2192,to study the characteristics of these clus-ters using UBV I CCD photometry.These clusters are lo-cated in the direction of anti-galactic centre.To date there is published only one photometric study of these clusters,which was given by Phelps et al.(1994)who presented non-calibrated BV CCD photometry of these clusters.From the instrumental color-magnitude diagrams of these clus-ters Phelps et al.estimated the ages of these clusters using the morphological age indicators,obtaining the values of 1.5Gyrs for NGC 1798and 1.1Gyrs for NGC 2192.However,no other properties of these clusters are yet known.In this paper we present a study of NGC 1798and NGC 2192based on UBV I CCD photometry.We have es-timated the basic parameters of these clusters:size,red-dening,metallicity,distance,and age.Also we have derived the luminosity function of the main sequence stars in these clusters.Section 2describes the observations and data re-duction.Sections 3and 4present the analysis for NGC 1798c 1998RAS2Hong Soo Park and Myung Gyoon Leeand NGC2192,respectively.Section5discusses the results. Finally Section6summarizes the primary results.2OBSER V ATIONS AND DATA REDUCTION 2.1ObservationsUBV I CCD images of NGC1798and NGC2192were obtained using the Photometrics512CCD camera at the Sobaeksan Observatory61cm telescope in Korea for several observing runs between1996November and1997October. We have used also BV CCD images of the central region of NGC1798obtained by Chul Hee Kim using the Tek1024 CCD camera at the Vainu Bappu Observatory2.3m tele-scope in India on March4,1998.The observing log is given in Table1.The original CCD images wereﬂattened after bias sub-traction and several exposures for eachﬁlter were combined into a single image for further reduction.The sizes of the ﬁeld in a CCD image are4′.3×4′.3for the PM512CCD im-age,and10′.6×10′.6for the Tek1024CCD image.The gain and readout noise are,respectively,9electrons/ADU and 10.4electrons for the PM512CCD,and9electrons/ADU and10.4electrons for the Tek1024CCD.Figs.1and2illustrate grey scale maps of the V CCD images of NGC1798and NGC2192made by mosaicing the images of the observed regions.It is seen from theseﬁgures that NGC1798is a relatively rich open cluster,while NGC 2192is a relatively poor open cluster.2.2Data ReductionInstrumental magnitudes of the stars in the CCD images were obtained using the digital stellar photometry reduction program IRAF⋆/DAOPHOT(Stetson1987,Davis1994). The resulting instrumental magnitudes were transformed onto the standard system using the standard stars from Lan-dolt(1992)and the M67stars in Montgomery et al.(1993) observed on the same photometric nights.The transforma-tion equations areV=v+a V(b−v)+k V X+Z V,(B−V)=a BV(b−v)+k BV X+Z BV,(U−B)=a UB(u−b)+k UB X+Z UB,and(V−I)=a V I(v−i)+k V I X+Z V I,where the lower case symbols represent instrumental mag-nitudes derived from the CCD images and the upper case symbols represent the standard system values.X is the air-mass at the midpoint of the observations.The results of the transformation are summarized in Table2.The data ob-tained on non-photometric nights were calibrated using the photometric data for the overlapped region.The total number of the measured stars is1,416for NGC1798and409for NGC2192.Tables3and4list the⋆IRAF is distributed by the National Optical Astronomy Ob-servatories,which are operated by the Association of Universities for Research in Astronomy,Inc.under contract with the National Science Foundation.photometry of the bright stars in the C-regions of NGC1798 and NGC2192,respectively.The X and Y coordinates listed in Table3and4are given in units of CCD pixel(=0′′.50). The X and Y values are increasing toward north and west, respectively.We have divided the entire region of theﬁelds into sev-eral regions,as shown in Figs.1and2,for the analysis of the data.The C-region represents the central region of the cluster,and the F-regions(F,Fb,Fir,and Fi regions)rep-resent the controlﬁeld regions,and the N-region represents the intermediate region between the central region and the ﬁeld region.The radius of the C-region is300pixel for NGC 1798and NGC2192.The ratio of the areas of the C-region, N-region,Fb-region,and(Fi+Fir)-regions for NGC1798is 1:1.50:1.00:1.07,and the ratio of the areas of the C-region, N-region,and F-region for NGC2192is1:1.26:0.98.3ANALYSIS FOR NGC17983.1The Size of the ClusterWe have investigated the structure of NGC1798using star-counts.The centre of the cluster is estimated to be at the po-sition of(X=710pixel,Y=1110pixel),using the centroid method.Fig.3illustrates the projected surface number den-sity proﬁle derived from counting stars with V<19.5mag in the entire CCDﬁeld.The magnitude cutoﬀfor starcounts was set so that the counts should be free of any photometric incompleteness problem.Fig.3shows that most of the stars in NGC1798are concentrated within the radius of250pixel (=125′′),and that the outskirts of the cluster extend out to about500pixel(=250′′)from the center.The number density changes little with radius beyond500pixel,show-ing that the outer region of the observedﬁeld can be used as a controlﬁeld.Therefore we have estimated the approxi-mate size of NGC1798for which the cluster blends in with theﬁeld to be about500′′in diameter,which corresponds to a linear size of10.2pc for the distance of NGC1798as determined below.3.2Color-Magnitude DiagramsFigs.4and5show the V−(B−V)and V−(V−I)color-magnitude diagrams(CMDs)of the measured stars in the observed regions in NGC1798.Theseﬁgures show that the C-region consists mostly of the members of NGC1798with some contamination of theﬁeld stars,while the F-regions consist mostly of theﬁeld stars.The N-region is intermediate between the C-region and the F-region.The distinguishable features seen in the color-magnitude diagrams of the C-region are:(a)There is a well-deﬁned main sequence the top of which is located at V≈16 mag;(b)There is seen a distinct gap at V≈16.2mag in the main sequence,which is often seen in other old open clusters (e.g.M67);(c)There is a poorly deﬁned red giant branch and these is seen some excess of stars around(B−V)=1.3 and V=15.6mag on this giant branch,which is remarked by the small box in theﬁgures.This may be a random excess of stars.However,the positions of the stars in the CMDs are consistent with the positions of known red giant clump in other old open clusters.Therefore most of these stars arec 1998RAS,MNRAS000,1–8UBV I CCD photometry of two old open clusters NGC1798and NGC21923probably red giant clump stars;and(d)There are a small number of stars along the locus of the red giant branch. 3.3Reddening and MetallicityNGC1798is located close to the galactic plane in the anti-galactic centre direction(b=4◦.85and l=160◦.76)so that it is expected that the reddening toward this cluster is signiﬁcant.We have estimated the reddening for NGC1798 using two methods as follows.First we have used the mean color of the red giant clump.Janes&Phelps(1994)estimated the mean color and magnitude of the red giant clump in old open clusters to be (B−V)RGC=0.87±0.02and M V,RGC=0.59±0.09,when the diﬀerence between the red giant clump and the main sequence turn-oﬀof the clusters,δV,is smaller than one. The mean color of the red giant clump in the C-region is estimated to be(B−V)RGC=1.34±0.01((V−I)RGC= 1.47±0.01,and(U−B)RGC=1.62±0.04),and the cor-responding mean magnitude is V RGC=15.57±0.05.δV is estimated to be0.8±0.2,which is the same value derived by Phelps et al.(1994).From these data we have derived a value of the reddening,E(B−V)=0.47±0.02.Secondly we have used the color-color diagram to es-timate the reddening and the metallicity simultaneously. We haveﬁtted the mean colors of the stars in the C-region with the color-color relation used in the Padova isochrones (Bertelli et al.1994).This process requires iteration,because we need to know the age of the cluster as well as the red-dening and metallicity.We have iterated this process until all three parameters are stabilized.Fig.6illustrates the results ofﬁtting in the(U−B)−(B−V)color-color diagram.It is shown in thisﬁgure that the stars in NGC1798are reasonablyﬁtted by the color-color relation of the isochrones for[Fe/H]=−0.47±0.15 with a reddening value of E(B−V)=0.55±0.05. The error for the metallicity,0.15,was estimated by com-paring isochrones with diﬀerent metallicities.As a reference the mean locus of the giants for solar abundance given by Schmidt-Kaler(1982)is also plotted in Fig.6.Finally we derive a mean value of the two estimates for the reddening, E(B−V)=0.51±0.04.3.4DistanceWe have estimated the distance to NGC1798using two methods as follows.First we have used the mean magnitude of the red giant clump.We have derived a value of the ap-parent distance modulus(m−M)V=14.98±0.10from the values for the mean magnitudes of the red giant clump stars described above.Secondly we have used the the zero-age main sequence (ZAMS)ﬁtting,following the method described in Vanden-Berg&Poll(1989).VandenBerg&Poll(1989)presented the semi-empirical ZAMS as a function of the metallicity[Fe/H] and the helium abundance Y:V=M V(B−V)+δM V(Y)+δM V([Fe/H])whereδM V(Y)=2.6(Y−0.27)andδM V([Fe/H])=−[Fe/H](1.444+0.362[Fe/H]).Before the ZAMSﬁtting,we subtracted statistically the contribution due to theﬁeld stars in the CMDs of the C-region using the CMDs of the Fb-region for BV photometry and the CMDs of the Fi+Fir region for V I photometry. The size of the bin used for the subtraction is∆V=0.25 and∆(B−V)=0.1.The resulting CMDs are displayed in Fig.7.We used the metallicity of[Fe/H]=–0.47as derived above and adopted Y=0.28which is the mean value for old open clusters(Gratton1982).Using this method we have obtained a value of the apparent distance modulus(m−M)V=14.5±0.2.Finally we calculate a mean value of the two estimates,(m−M)V=14.7±0.2.Adopting the extinction law of A V=3.2E(B−V),we derive a value of the intrinsic distance modulus(m−M)0=13.1±0.2.This corresponds to a distance of d=4.2±0.3kpc.3.5AgeWe have estimated the age of NGC1798using two methods as follows.First we have used the morphological age index (MAI)as described in Phelps et al.(1994).Phelps et al. (1994)and Janes&Phelps(1994)presented the MAI–δV relation,MAI[Gyrs]=0.73×10(0.256δV+0.0662δV2).From the value ofδV derived above,0.8±0.2mag,we obtain a value for the age,MAI=1.3±0.2Gyrs.Secondly we have estimated the age of the cluster us-ing the theoretical isochrones given by the Padova group (Bertelli et al.1994).Fitting the isochrones for[Fe/H]=–0.47to the CMDs of NGC1798,as shown in Fig.8,we estimate the age to be1.4±0.2Gyrs.Both results agree very well.3.6Luminosity FunctionWe have derived the V luminosity functions of the main sequence stars in NGC1798,which are displayed in Fig.9.The Fb-region was used for subtraction of theﬁeld star contribution from the C-region and the magnitude bin size used is0.5mag.This controlﬁeld may not be far enough from the cluster to derive theﬁeld star contribution.If so,we might have oversubtracted theﬁeld contribution,obtaining ﬂatter luminosity functions than true luminosity functions. However,the fraction of the cluster members in thisﬁeld must be,if any,very low,because the surface number density of this region is almost constant with the radius as shown in Fig.3.The luminosity function of the C-region in Fig. 9(a)increases rapidly up to V≈16.5mag,and stays almost ﬂat for V>16.5mag.The luminosity functions of the N-region and the(R+Fir)-region are steeper than that of the C-region.A remarkable drop is seen at V=16.2mag(M V= 1.5mag)in the luminosity function of the C-region based on smaller bin size of0.2mag in Fig.9(b).This corresponds to the main sequence gap described above.4ANALYSIS FOR NGC21924.1The Size of the ClusterWe have investigated the structure of NGC2192using star-counts.We could not use the centroid method to estimatec 1998RAS,MNRAS000,1–84Hong Soo Park and Myung Gyoon Leethe centre of this cluster,because this cluster is too sparse. So we have used eye-estimate to determine the centre of the cluster to be at the position of(X=465pixel,Y=930 pixel).Fig.10illustrates the projected surface number den-sity proﬁle derived from counting stars with V<18mag in the entire CCDﬁeld.The magnitude cutoﬀfor starcounts was set so that the counts should be free of any photomet-ric incompleteness problem.Fig.10shows that most of the stars in NGC2192are concentrated within the radius of200 pixel(=100′′),and that the outskirts of the cluster extend out to about440pixel(=220′′)from the centre.Therefore the approximate size of NGC2192is estimated to be about 440′′in diameter,which corresponds to a linear size of7.5pc for the distance of NGC2192as determined below.4.2Color-Magnitude DiagramsFigs.11and12show the V−(B−V)and V−(V−I)color-magnitude diagrams of the measured stars in the observed regions in NGC2192.The distinguishable features seen in the color-magnitude diagrams of the C-region are:(a)There is a well-deﬁned main sequence the top of which is located at V≈14mag;(b)There are a group of red giant clump stars at(B−V)=1.1and V=14.2mag,which are remarked by the small box in theﬁgures;and(c)There are a small number of stars along the locus of the red giant branch. 4.3Reddening and MetallicityNGC2192is located11degrees above the galactic plane in the anti-galactic centre direction(b=10◦.64and l= 173◦.41)but higher than NGC1798so that it is expected that the reddening toward this cluster is signiﬁcant but smaller than that of NGC1798.We have estimated the red-dening for NGC2192using two methods as applied for NGC 1798.First we have used the mean color of the red giant clump.The mean color of the red giant clump in the C-region is estimated to be(B−V)RGC=1.08±0.01((V−I)RGC=1.07±0.01,and(U−B)RGC=0.61±0.02),and the corresponding mean magnitude is V RGC=14.20±0.05.δV is estimated to be0.6±0.2,which is similar to the value derived by Phelps et al.(1994).From these data we have derived a value of the reddening,E(B−V)=0.19±0.03.Secondly we have used the color-color diagram to es-timate the reddening and the metallicity simultaneously. We haveﬁtted the mean colors of the stars in the C-region with the color-color relation used in the Padova isochrones (Bertelli et al.1994).Fig.13illustrates the results ofﬁt-ting in the(U−B)−(B−V)color-color diagram.It is shown in thisﬁgure that the stars in NGC2192are rea-sonablyﬁtted by the color-color relation of the isochrones for[Fe/H]=−0.31±0.15dex with a reddening value of E(B−V)=0.21±0.01.The error for the metallicity,0.15, was estimated by comparing isochrones with diﬀerent metal-licities.As a reference the mean locus of the giant for solar abundance given by Schmidt-Kaler is also plotted in Fig.13. Finally we derive a mean value of the two estimates for the reddening,E(B−V)=0.20±0.03.4.4DistanceWe have estimated the distance to NGC2192using two methods as for NGC1798.First we have used the mean magnitude of the red giant clump.We have derived a value of the apparent distance modulus(m−M)V=13.61±0.10 from the values for the mean magnitudes of the red giant clump stars described previously.Secondly we have used the ZAMSﬁtting.Before the ZAMSﬁtting,we subtracted statistically the contribution due to theﬁeld stars in the CMDs of the C-region using the CMDs of the F-region.The size of the bin used for the subtraction is∆V=0.25and∆(B−V)=0.1.The resulting CMDs are displayed in Fig.14.We used the metallicity of[Fe/H]=–0.31as derived before and adopted Y=ing this method we have obtaineda value of the apparent distance modulus(m−M)V=13.1±0.2.Finally we calculate a mean value of the two estimates,(m−M)V=13.3±0.2.Adopting the extinction law of A V=3.2E(B−V),we derive a value of the intrinsic distance modulus(m−M)0=12.7±0.2.This corresponds to a distance of d=3.5±0.3kpc.4.5AgeWe have estimated the age of NGC2192using two methods as follows.First we have used the morphological age index. From the value ofδV derived above,0.6±0.2mag,we obtain a value for the age,MAI=1.1±0.2Gyrs.Secondly we have estimated the age of the cluster using the theoretical isochrones given by the Padova group(Bertelli et al.1994). Fitting the isochrones for[Fe/H]=–0.31to the CMDs of NGC2192,as shown in Fig.15,we estimate the age to be 1.1±0.1Gyrs.Both results agree very well.4.6Luminosity FunctionWe have derived the V luminosity functions of the main sequence stars in NGC2192,which are displayed in Fig.16.The F-region was used for subtraction of theﬁeld star contribution from the C-region.The luminosity function of the C-region in Fig.16(a)increases rapidly up to V≈14 mag,and stays almostﬂat for V>15mag.The luminosity function of the N-region is steeper than that of the C-region. Fig.16(b)displays a comparison of the luminosity functions of NGC1798,NGC2192,and NGC7789which is another old open cluster of similar age(Roger et al.1994).Fig.16(b) shows that the luminosity functions of these clusters are similar in that they are almostﬂat in the faint part.The ﬂattening of the faint part of the luminosity functions of old open clusters has been known since long,and is believed to be due to evaporation of low mass stars(Friel1995).5DISCUSSIONWe have determined the metallicity and distance of NGC 1798and NGC2192in this study.We compare them with those of other old open clusters here.Fig.17illustrates the radial metallicity gradient of the old open clusters com-piled by Friel(1995)and supplemented by the data in Wee&Lee(1996)and Lee(1997).Fig.17shows that thec 1998RAS,MNRAS000,1–8UBV I CCD photometry of two old open clusters NGC1798and NGC21925 mean metallicity decreases as the galactocentric distance in-creases.The positions of NGC1798and NGC2192we haveobtained in this study are consistent with the mean trendof the other old open clusters.The slope we have deter-mined for the entire sample including these two clusters is∆[Fe/H]/R GC=−0.086±0.011dex/kpc,very similar tothat given in Friel(1995),∆[Fe/H]/R GC=−0.091±0.014dex/kpc.There are only four old open clusters located beyondR GC=13kpc in Fig.17.These four clusters follow the meantrend of decreasing outward.However,the number of theclusters is not large enough to decide whether the metallictykeeps decreasing outward or it stops decreasing somewherebeyond R GC=13kpc and stays constant.Further studiesof more old open clusters beyond R GC=13kpc are neededto investigate this point.6SUMMARY AND CONCLUSIONSWe have presented UBV I photometry of old open clustersNGC1798and NGC2192.From the photometric data wehave determined the size,reddening,metallicity,distance,and age of these clusters.The luminosity functions of themain sequence stars in these clusters are similar to those ofthe other old open clusters.The basic properties of theseclusters we have determined in this study are summarizedin Table5.ACKNOWLEDGMENTSProf.Chul Hee Kim is thanked for providing the BV CCDimages of NGC1798.This research is supported in part bythe Korea Science and Engineering Foundation Grant No.95-0702-01-01-3.REFERENCESBertelli,G.,Bressan,A.,Chiosi,C.,Fagotto,F.,&Nasi,E.1994,A&AS,106,275Davis,L.E.,1994,A Reference Guide to the IRAF/DAOPHOTPackageFriel,E.D.1995,ARA&A,33,381Gratton,R.G.,1982,ApJ,257,640Janes,K.,&Phelps,R.L.1994,AJ,108,1773Landolt,A.U.,1992,AJ,104,340Lee,M.G.,1997,113,729Montgomery,K.A.,Marschall,L.A.,&Janes,K.A.1993,AJ,106,181Phelps,R.L.,Janes,K.,&Montgomery,K.A.,1994,AJ,107,1079Roger,C.M.,Paez,E.,Castellani,V.,&Staniero,O.1994,A&A,290,62Schmidt-Kaler,T.1982,in Landolt-Bornstein VI,2b(Berlin:Springer)Stetson,P.B.,1987,PASP,99,191VandenBerg,D.A.,&Poll,H.E.,1989,AJ,98,1451Wee,S.O.,&Lee,M.G.,1996,Jour.Korean Astro.Soc.,29,181This paper has been produced using the Royal AstronomicalSociety/Blackwell Science L A T E X styleﬁle.c 1998RAS,MNRAS000,1–86Hong Soo Park and Myung Gyoon LeeTable2.Transformation coeﬃcients for the standard stars.Date Color a k Z rms n(stars) 96.11.11V0.003–0.101–6.0400.00919(B−V) 1.090–0.155–0.4670.01317(U−B) 1.008–0.154–1.7110.03416 97.02.11V0.028–0.198–6.0500.00934(B−V) 1.150–0.118–0.5930.01035(U−B) 1.079–0.324–1.5750.03226(V−I)0.983–0.1300.3020.00829 97.02.13V–0.007–0.176–6.0050.01034(B−V) 1.160–0.133–0.5880.01334(U−B) 1.079–0.349–1.5260.02320(V−I)0.986–0.0970.2710.01532 97.03.17V-0.019–0.225–6.0680.01548(B−V) 1.221–0.093–0.7810.01844(U−B) 1.008–0.309–1.4580.02735(V−I)0.956–0.1570.3140.02052Table 3.UBV I photometry of the bright stars in the C-regionof NGC1798.ID X[px]Y[px]V(B−V)(U−B)(V−I) 1723.21110.815.8600.7790.3240.9143680.51092.015.380 1.3250.627 1.5254746.01160.315.9050.7410.2200.9545741.81165.514.380 1.4830.880 1.6336684.61030.115.8140.8640.342 1.0427745.31073.315.3760.9200.237 1.0849707.21038.814.777 1.4860.723 1.653 11549.41153.715.305 1.5290.915 1.666 12547.91162.714.846 1.6270.947 1.764 13602.31068.015.325 1.5240.847 1.660 14598.11080.815.776 1.3340.575 1.490 15525.81116.815.510 1.0020.346 1.169 18833.91144.214.976 1.4940.866 1.578 21981.91125.114.711 1.510 1.031 1.440 26835.71098.015.706 1.3180.468 1.407 27656.51022.115.653 1.4160.180 1.571 28659.8955.915.216 1.2160.384 1.385 29656.4976.813.6670.6610.3100.777 30674.2822.512.0330.4450.0500.532 33668.7966.315.687 1.3130.641 1.446 34764.3879.915.460 1.3370.508 1.408 38749.01008.115.359 1.3220.744 1.426 39774.4965.915.545 1.3040.590 1.402 43571.8930.612.875 1.970 1.264 2.252 45512.8908.715.810 1.3480.474 1.485 46642.1949.515.622 1.3290.590 1.477 48732.61211.315.821 1.3530.643 1.492 50678.81210.515.297 1.3680.770 1.501 51742.11244.414.7990.747-0.046 1.216 52673.11250.215.180 1.3060.780 1.457 56728.21183.415.665 1.4050.803 1.545 57614.91229.413.085 2.1770.635 3.581 62563.01338.415.9090.7060.6380.778 65913.91203.013.686 1.2490.851 1.354 67792.71314.215.6010.9750.470 1.275 68932.61268.414.4660.7170.1970.763 70801.11338.715.561 1.3270.785 1.442 72827.61315.513.5260.9670.521 1.161 78444.01094.814.605 1.5790.825 1.723 82684.3822.615.8280.9510.128 1.189 97598.51330.315.580 1.3360.710 1.483Table4.UBV I photometry of the bright stars in the C-region of NGC2192.ID X[px]Y[px]V(B−V)(U−B)(V−I) 4502.4879.614.9830.4320.1140.424 5537.5891.712.9690.9240.4750.930 7535.1931.315.1700.598-0.2010.483 9434.8947.613.6760.570-0.0560.648 10562.3991.914.037 1.0960.612 1.043 11412.9884.814.165 1.0780.608 1.072 12598.4778.715.1840.464-0.0300.477 13595.9930.014.4800.5020.1720.549 14625.8971.814.7430.7200.0890.741 16614.3684.214.3970.7660.1750.758 17463.3689.514.7660.4660.0120.480 18553.3710.715.4920.447-0.1350.428 19591.7723.814.4660.5450.0930.566 21623.0867.712.634 1.1020.713 1.108 22648.0929.015.3380.3940.1150.399 26461.1711.314.162 1.0860.670 1.091 29583.9802.315.0940.600-0.0160.644 32720.4889.515.3260.3700.0710.410 37395.9778.315.2640.3960.0310.379 38405.5778.514.8900.4810.1010.409 40258.21042.915.4310.4240.1200.484 42376.6883.915.1870.4550.1360.470 44180.1908.714.3200.2870.1120.295 47367.1938.314.8970.4030.2520.459 51183.3998.714.202 1.1140.552 1.108 55475.81052.514.6370.3880.1670.465 56389.11055.113.6060.6520.2560.777 58481.41074.114.3470.4500.1150.520 59514.91085.115.0970.4290.1370.517 60478.51089.614.2360.5280.1590.582 61498.51127.914.6090.5090.1350.578 63278.71154.414.372 1.0700.599 1.053 64527.41158.715.4180.4460.2450.510 65343.41167.814.266 1.0120.618 1.035 66362.21169.314.2660.8790.4690.879 74346.91039.814.5810.4520.1240.510 77688.91124.215.4200.3900.2900.400 83678.91079.715.4950.4640.1070.504 84624.21042.813.5010.3850.2170.389 88579.61028.914.8250.4560.1620.530 Table5.Basic properties of NGC1798and NGC2192.Parameter NGC1798NGC2192RA(2000)5h11m40s6h15m11sDEC(2000)47◦40′37′′39◦51′1′′l160◦.76173◦.41b4◦.8510◦.64Age1.4±0.2Gyrs1.1±0.1GyrsE(B−V)0.51±0.040.20±0.04[Fe/H]−0.47±0.15dex−0.31±0.15dex(m−M)013.1±0.212.7±0.2distance4.2±0.3kpc3.5±0.3kpcR GC12.5kpc11.9kpcz355pc646pcdiameter10.2pc(8′.3)7.5pc(7′.3)c 1998RAS,MNRAS000,1–8UBV I CCD photometry of two old open clusters NGC1798and NGC21927 Table1.Observing log for NGC1798and NGC2192.Date Target Filter Seeing Telescope Condition96.11.11NCC1798UBV2′′.3SAO a-61cm Photometric97.01.12NCC1798UBV I2′′.2SAO-61cm Non-photometric97.02.12NCC1798,NGC2192UBV I3′′.3SAO-61cm Photometric97.02.13NCC1798,NGC2192UBV I2′′.5SAO-61cm Non-photometric97.03.17NCC2192UBV I2′′.2SAO-61cm Photometric97.10.21NCC1798BV I2′′.3SAO-61cm Non-photometric97.10.24NCC2192BV I2′′.7SAO-61cm Non-photometric97.03.04NCC1798BV2′′.7VBO b-2.3m Non-photometrica Sobaeksan Astronomical Observatoryb Vainu Bappu Observatoryc 1998RAS,MNRAS000,1–8。

中国天文学会天文学名词审定委员会第1-6批天文学名词的推荐译名

中国天⽂学会天⽂学名词审定委员会第1-6批天⽂学名词的推荐译名The 1st - 6th Drafts for the Chinese-Translation of Astronomical Terms recommanded byThe Astronomical Terminology Committee of the CASabsolute stability 绝对稳定性absorbing dust mass 致吸尘物质absorption trough 吸收槽abundance standard 丰度标准星accreting binary 吸积双星accretion column 吸积柱accretion flow 吸积流accretion mound 吸积堆accretion ring 吸积环accretion stream 吸积流acoustic mode 声模active binary 活动双星active chromosphere binary 活动⾊球双星active chromosphere star 活动⾊球星active optics 主动光学actuator 促动器Adams ring 亚当斯环adaptive optics ⾃适应光学additional perturbation 附加摄动AGB, asymptotic giant branch 渐近巨星⽀Alexis, Array of Low-Energy X-ray 〈阿列克希斯〉低能 X 射线Imaging Sensors 成象飞⾏器AM Herculis star 武仙 AM 型星amplitude spectrum 变幅谱angular elongation 距⾓anonymous galaxy 未名星系anonymous object 未名天体anti-jovian point 对⽊点annular-total eclipse 全环⾷aperture photometry 孔径测光APM, Automated Photographic Measuring 〈APM〉底⽚⾃动测量仪systemapoapse 远质⼼点apoapse distance 远质⼼距apogalacticon 远银⼼点apomartian 远⽕点apparent association 表观成协apparent luminosity function 视光度函数apparent superluminal motion 视超光速运动apsidal advance 拱线进动apsidal precession 拱线进动Arcturus group ⼤⾓星群area image sensor ⾯成象敏感器area photometry ⾯源测光area spectroscopy ⾯源分光argument of pericentre 近⼼点幅⾓ASCA, Advanced Satellite for Cosmology 〈ASCA〉宇宙学和天体物理学and Astrophysics ⾼新卫星asteroidal dynamics ⼩⾏星动⼒学asteroidal resonance ⼩⾏星共振asteroid family ⼩⾏星族asteroid-like object 类⼩⾏星天体asteroseismology 星震学astration 物质改造astroparticle physics 天⽂粒⼦物理学astrostatistics 天⽂统计学asymptotic branch 渐近⽀asymptotic branch giant 渐近⽀巨星atmospheric parameter ⼤⽓参数ATNT, Australia Telescope National 澳⼤利亚国⽴望远镜FacilityATT, Advanced Technology Telescope 〈ATT〉⾼新技术望远镜automated measuring machine 天⽂底⽚⾃动测量仪automatic photooelectric telescope ⾃动光电测光望远镜( APT )AXAF, Advanced X-ray Astrophysical ⾼新X射线天体物理台FacilityBaade's window 巴德窗Baade—Wesselink analysis 巴德—韦塞林克分析Baade—Wesselink mass 巴德—韦塞林克质量Baade—Wesselink method 巴德—韦塞林克⽅法Baade—Wesselink radius 巴德—韦塞林克半径background galaxy 背景星系Barnard's galaxy ( NGC 6822 ) 巴纳德星系barycentric dynamical time ( TDB ) 质⼼⼒学时Belinda 天卫⼗四Bianca 天卫⼋bidimensional spectrography ⼆维摄谱bidimensional spectroscopy ⼆维分光Big-Bang nucleosynthesis ⼤爆炸核合成binarity 成双性binary asteroid 双⼩⾏星binary flare star 耀变双星binary millisecond pulsar 毫秒脉冲双星binary protostar 原双星bioastronomy ⽣物天⽂学bipolar jet 双极喷流bipolar outflow 偶极外向流bipolar planetary nebula 双极⾏星状星云blazar 耀变体blazarlike activity 类耀活动blazarlike object 耀变体Black-eye galaxy ( M 64 ) ⿊眼星系BL Lacertae object 蝎虎天体BL Lacertid 蝎虎天体blue compact galaxy ( BCG ) 蓝致密星系blue straggler 蓝离散星bolometric albedo 热反照率bolometric light curve 全波光变曲线bolometric temperature 热温度Bootes void 牧夫巨洞bow-shock nebula ⼸形激波星云box photometry ⽅格测光broad-band imaging 宽波段成象broad-line radio galaxy ( BLRG ) 宽线射电星系buried channel CCD 埋沟型 CCDButterfly nebula 蝴蝶星云BY Draconis star 天龙 BY 型星BY Draconis variable 天龙 BY 型变星CAMC, Carlsberg Automatic Meridian 卡尔斯伯格⾃动⼦午环Circlecannibalism 吞⾷cannibalized galaxy 被吞星系cannibalizing galaxy 吞⾷星系cannibalizing of galaxies 星系吞⾷carbon dwarf 碳矮星Cassegrain spectrograph 卡焦摄谱仪Cassini 〈卡西尼〉⼟星探测器Cat's Eye nebula ( NGC 6543 ) 猫眼星云CCD astronomy CCD 天⽂学CCD camera CCD 照相机CCD photometry CCD 测光CCD spectrograph CCD 摄谱仪CCD spectrum CCD 光谱celestial clock 天体钟celestial mechanician 天体⼒学家celestial thermal background 天空热背景辐射celestial thermal background radiation 天空热背景辐射central overlap technique 中⼼重迭法Centaurus arm 半⼈马臂Cepheid distance 造⽗距离CFHT, Canada-Franch-Hawaii Telecope 〈CFHT〉望远镜CGRO, Compton Gamma-Ray Observatory 〈康普顿〉γ射线天⽂台chaos 混沌chaotic dynamics 混沌动⼒学chaotic layer 混沌层chaotic region 混沌区chemically peculiar star 化学特殊星Christmas Tree cluster ( NGC 2264 ) 圣诞树星团chromosphere-corona transition zone ⾊球-⽇冕过渡层chromospheric activity ⾊球活动chromospherically active banary ⾊球活动双星chromospherically active star ⾊球活动星chromospheric line ⾊球谱线chromospheric matirial ⾊球物质chromospheric spectrum ⾊球光谱CID, charge injected device CID、电荷注⼊器件circular solution 圆轨解circumnuclear star-formation 核周产星circumscribed halo 外接⽇晕circumstellar dust disk 星周尘盘circumstellar material 星周物质circumsystem material 双星周物质classical Algol system 经典⼤陵双星classical quasar 经典类星体classical R Coronae Borealis star 经典北冕 R 型星classical T Tauri star 经典⾦⽜ T 型星Clementine 〈克莱芒蒂娜〉环⽉测绘飞⾏器closure phase imaging 锁相成象cluster centre 团中⼼cluster galaxy 团星系COBE, Cosmic Background Explorer 宇宙背景探测器coded mask imaging 编码掩模成象coded mask telescope 编码掩模望远镜collapsing cloud 坍缩云cometary burst 彗暴cometary dynamics 彗星动⼒学cometary flare 彗耀cometary H Ⅱ region 彗状电离氢区cometary outburst 彗爆发cometary proplyd 彗状原⾏星盘comet shower 彗星⾬common proper-motion binary 共⾃⾏双星common proper-motion pair 共⾃⾏星对compact binary galaxy 致密双重星系compact cluster 致密星团; 致密星系团compact flare 致密耀斑composite diagram method 复合图法composite spectrum binary 复谱双星computational astrophysics 计算天体物理computational celestial mechanics 计算天体⼒学contact copying 接触复制contraction age 收缩年龄convective envelope 对流包层cooling flow 冷却流co-orbital satellite 共轨卫星coplanar orbits 共⾯轨道Copernicus 〈哥⽩尼〉卫星coprocessor 协处理器Cordelia 天卫六core-dominated quasar ( CDQ ) 核占优类星体coronal abundance 冕区丰度coronal activity 星冕活动、⽇冕活动coronal dividing line 冕区分界线coronal gas 星冕⽓体、⽇冕⽓体coronal green line 星冕绿线、⽇冕绿线coronal helmet 冕盔coronal magnetic energy 冕区磁能coronal red line 星冕红线、⽇冕红线cosmic abundance 宇宙丰度cosmic string 宇宙弦cosmic void 宇宙巨洞COSMOS 〈COSMOS〉底⽚⾃动测量仪C-O white dwarf 碳氧⽩矮星Cowling approximation 柯林近似Cowling mechnism 柯林机制Crescent nebula ( NGC 6888 ) 蛾眉⽉星云Cressida 天卫九critical equipotential lobe 临界等位瓣cross-correlation method 交叉相关法cross-correlation technique 交叉相关法cross disperser prism 横向⾊散棱镜crustal dynamics 星壳动⼒学cryogenic camera 致冷照相机cushion distortion 枕形畸变cut-off error 截断误差Cyclops project 〈独眼神〉计划D abundance 氘丰度Dactyl 艾卫dark halo 暗晕data acquisition 数据采集decline phase 下降阶段deep-field observation 深天区观测density arm 密度臂density profile 密度轮廓dereddening 红化改正Desdemona 天卫⼗destabiliizing effect 去稳效应dew shield 露罩diagonal mirror 对⾓镜diagnostic diagram 诊断图differential reddening 较差红化diffuse density 漫射密度diffuse dwarf 弥漫矮星系diffuse X-ray 弥漫 X 射线diffusion approximation 扩散近似digital optical sky survey 数字光学巡天digital sky survey 数字巡天disappearance 掩始cisconnection event 断尾事件dish 碟形天线disk globular cluster 盘族球状星团dispersion measure 频散量度dissector 析象管distance estimator 估距关系distribution parameter 分布参数disturbed galaxy 受扰星系disturbing galaxy 扰动星系Dobsonian mounting 多布森装置Dobsonian reflector 多布森反射望远镜Dobsonian telescope 多布森望远镜dominant galaxy 主星系double-mode cepheid 双模造⽗变星double-mode pulsator 双模脉动星double-mode RR Lyrae star 双模天琴 RR 型星double-ring galaxy 双环星系DQ Herculis star 武仙 DQ 型星dredge-up 上翻drift scanning 漂移扫描driving system 驱动系统dumbbell radio galaxy 哑铃状射电星系Du Pont Telescope 杜邦望远镜dust ring 尘环dwarf carbon star 碳矮星dwarf spheroidal 矮球状星系dwarf spheroidal galaxy 矮球状星系dwarf spiral 矮旋涡星系dwarf spiral galaxy 矮旋涡星系dynamical age 动⼒学年龄dynamical astronomy 动⼒天⽂dynamical evolution 动⼒学演化Eagle nebula ( M 16 ) 鹰状星云earty cluster 早型星系团early earth 早期地球early planet 早期⾏星early-stage star 演化早期星early stellar evolution 恒星早期演化early sun 早期太阳earth-approaching asteroid 近地⼩⾏星earth-approaching comet 近地彗星earth-approaching object 近地天体earth-crossing asteroid 越地⼩⾏星earth-crossing comet 越地彗星earth-crossing object 越地天体earth orientation parameter 地球定向参数earth rotation parameter 地球⾃转参数eccentric-disk model 偏⼼盘模型effect of relaxation 弛豫效应Egg nebula ( AFGL 2688 ) 蛋状星云electronographic photometry 电⼦照相测光elemental abundance 元素丰度elliptical 椭圆星系elliptical dwarf 椭圆矮星系emulated data 仿真数据emulation 仿真encounter-type orbit 交会型轨道enhanced network 增强络equatorial rotational velocity ⾚道⾃转速度equatorium ⾏星定位仪equipartition of kinetic energy 动能均分eruptive period 爆发周期Eskimo nebula ( NGC 2392 ) 爱斯基摩星云estimated accuracy 估计精度estimation theory 估计理论EUVE, Extreme Ultraviolet Explorer 〈EUVE〉极紫外探测器Exclamation Mark galaxy 惊叹号星系Exosat 〈Exosat〉欧洲 X 射线天⽂卫星extended Kalman filter 扩充卡尔曼滤波器extragalactic jet 河外喷流extragalactic radio astronomy 河外射电天⽂extrasolar planet 太阳系外⾏星extrasolar planetary system 太阳系外⾏星系extraterrestrial intelligence 地外智慧⽣物extreme helium star 极端氦星Fabry-Perot imaging spectrograph 法布⾥-珀罗成象摄谱仪Fabry-Perot interferometry 法布⾥-珀罗⼲涉测量Fabry-Perot spectrograph 法布⾥-珀罗摄谱仪face-on galaxy 正向星系face-on spiral 正向旋涡星系facility seeing ⼈为视宁度fall 见落陨星fast pulsar 快转脉冲星fat zero 胖零Fermi normal coordinate system 费⽶标准坐标系Fermi-Walker transportation 费⽶-沃克移动fibre spectroscopy 光纤分光field centre 场中⼼field galaxy 场星系field pulsar 场脉冲星filter photography 滤光⽚照相观测filter wheel 滤光⽚转盘find 发见陨星finder chart 证认图finderscope 寻星镜first-ascent giant branch 初升巨星⽀first giant branch 初升巨星⽀flare puff 耀斑喷焰flat field 平场flat field correction 平场改正flat fielding 平场处理flat-spectrum radio quasar 平谱射电类星体flux standard 流量标准星flux-tube dynamics 磁流管动⼒学f-mode f 模、基本模following limb 东边缘、后随边缘foreground galaxy 前景星系foreground galaxy cluster 前景星系团formal accuracy 形式精度Foucaultgram 傅科检验图样Foucault knife-edge test 傅科⼑⼝检验fourth cosmic velocity 第四宇宙速度frame transfer 帧转移Fresnel lens 菲涅尔透镜fuzz 展云Galactic aggregate 银河星集Galactic astronomy 银河系天⽂Galactic bar 银河系棒galactic bar 星系棒galactic cannibalism 星系吞⾷galactic content 星系成分galactic merge 星系并合galactic pericentre 近银⼼点Galactocentric distance 银⼼距galaxy cluster 星系团Galle ring 伽勒环Galilean transformation 伽利略变换Galileo 〈伽利略〉⽊星探测器gas-dust complex ⽓尘复合体Genesis rock 创世岩Gemini Telescope ⼤型双⼦望远镜Geoalert, Geophysical Alert Broadcast 地球物理警报⼴播giant granulation 巨⽶粒组织giant granule 巨⽶粒giant radio pulse 巨射电脉冲Ginga 〈星系〉X 射线天⽂卫星Giotto 〈乔托〉空间探测器glassceramic 微晶玻璃glitch activity ⾃转突变活动global change 全球变化global sensitivity 全局灵敏度GMC, giant molecular cloud 巨分⼦云g-mode g 模、重⼒模gold spot ⾦斑病GONG, Global Oscillation Network 太阳全球振荡监测GroupGPS, global positioning system 全球定位系统Granat 〈⽯榴〉号天⽂卫星grand design spiral 宏象旋涡星系gravitational astronomy 引⼒天⽂gravitational lensing 引⼒透镜效应gravitational micro-lensing 微引⼒透镜效应great attractor 巨引源Great Dark Spot ⼤暗斑Great White Spot ⼤⽩斑grism 棱栅GRO, Gamma-Ray Observatory γ射线天⽂台guidscope 导星镜GW Virginis star 室⼥ GW 型星habitable planet 可居住⾏星Hakucho 〈天鹅〉X 射线天⽂卫星Hale Telescope 海尔望远镜halo dwarf 晕族矮星halo globular cluster 晕族球状星团Hanle effect 汉勒效应hard X-ray source 硬 X 射线源Hay spot 哈伊斑HEAO, High-Energy Astronomical 〈HEAO〉⾼能天⽂台Observatoryheavy-element star 重元素星heiligenschein 灵光Helene ⼟卫⼗⼆helicity 螺度heliocentric radial velocity ⽇⼼视向速度heliomagnetosphere ⽇球磁层helioseismology ⽇震学helium abundance 氦丰度helium main-sequence 氦主序helium-strong star 强氦线星helium white dwarf 氦⽩矮星Helix galaxy ( NGC 2685 ) 螺旋星系Herbig Ae star 赫⽐格 Ae 型星Herbig Be star 赫⽐格 Be 型星Herbig-Haro flow 赫⽐格-阿罗流Herbig-Haro shock wave 赫⽐格-阿罗激波hidden magnetic flux 隐磁流high-field pulsar 强磁场脉冲星highly polarized quasar ( HPQ ) ⾼偏振类星体high-mass X-ray binary ⼤质量 X 射线双星high-metallicity cluster ⾼⾦属度星团;⾼⾦属度星系团high-resolution spectrograph ⾼分辨摄谱仪high-resolution spectroscopy ⾼分辨分光high - z ⼤红移Hinotori 〈⽕鸟〉太阳探测器Hipparcos, High Precision Parallax 〈依巴⾕〉卫星Collecting SatelliteHipparcos and Tycho Catalogues 〈依巴⾕〉和〈第⾕〉星表holographic grating 全息光栅Hooker Telescope 胡克望远镜host galaxy 寄主星系hot R Coronae Borealis star ⾼温北冕 R 型星HST, Hubble Space Telescope 哈勃空间望远镜Hubble age 哈勃年龄Hubble distance 哈勃距离Hubble parameter 哈勃参数。

IRAC and MIPS Observations of the Interacting Galaxies IC 2163 and NGC 2207 Clumpy Emission

a rXiv:as tr o-ph/65524v12May26Spitzer Space Telescope IRAC and MIPS Observations of the Interacting Galaxies IC 2163and NGC 2207:Clumpy Emission Debra Meloy Elmegreen 1,Bruce G.Elmegreen 2,Michele Kaufman 3,Kartik Sheth 4,Curtis Struck 5,Magnus Thomasson 6,and Elias Brinks 7ABSTRACT IC 2163and NGC 2207are interacting galaxies that have been well studied at optical and radio wavelengths and simulated in numerical models to repro-duce the observed kinematics and morphological features.Spitzer IRAC and MIPS observations reported here show over 200bright clumps from young star complexes.The brightest IR clump is a morphologically peculiar region of star formation in the western arm of NGC 2207.This clump,which dominates the H αand radio continuum emission from both galaxies,accounts for ∼12%of the total 24µm ﬂux.Nearly half of the clumps are regularly spaced along some ﬁlamentary structure,whether in the starburst oval of IC 2163or in the thin spiral arms of NGC 2207.This regularity appears to inﬂuence the clump luminosity function,making it peaked at a value nearly a factor of 10above the completeness limit,particularly in the starburst oval.This is unlike the optical clusters inside the clumps,which have a luminosity function consistent with the usual power law form.The giant IR clumps presumably formed by gravitational instabilities in the compressed gas of the oval and the spiral arms,whereas the individual clus-ters formed by more chaotic processes,such as turbulence compression,inside these larger-scale structures.Subject headings:galaxies:interacting—galaxies:individual(NGC2207)—galaxies:individual(IC2163)—infrared:galaxies1.IntroductionIC2163and NGC2207are interacting galaxies at a distance of35Mpc(Elmegreen et al. 1995a,hereafter Paper I).IC2163has an unusual bright oval of star formation shaped like an eyelid(“ocular”structure)at mid-radius that was predicted to be the result of large-scale gaseous shocks in simulations of grazing prograde encounters(Elmegreen et al.1991;Sundin 1993).Several ocular or post-ocular galaxies and their companions have been observed and modeled in detail by our group(Kaufman et al.1997;Kaufman et al.1999;Kaufman et al. 2002).The IC2163/NGC2207pair was observed at optical(Elmegreen et al.2000,hereafter Paper III;2001,hereafter Paper IV),millimeter(Thomasson2004),and radio wavelengths (Paper I;Elmegreen et al.1995b,hereafter Paper II).HST WFPC2observations were used to measure the properties of compact star clusters,many of which satisfy the criteria for super star clusters(SSC;Paper IV),and to determine the opacities of dust clouds in the foreground screen provided by NGC2207(Paper IV).The HST observations also showed that the clusters and stellar groupings inﬁve major star-forming regions are hierarchically distributed and that the fractal dimensions of the starﬁelds are comparable to the fractal dimension of interstellar gas(Elmegreen&Elmegreen2001).Also,HST PC images of the central region of NGC2207revealed what appears to be acoustic turbulence(Elmegreen et al.1998;Montenegro,Yuan&Elmegreen1999;for other examples,see Martini et al.2003).Detailed numerical simulations have reproduced over two dozen dynamical and struc-tural details of the IC2163/NGC2207pair(Paper II;Paper III;Struck et al.2005,hereafter Paper V).The simulations suggest that NGC2207is moving prograde in the plane of IC 2163,causing tidal forces that initially compressed and stretched IC2163so that the stars were jarred into caustic patterns and the gas was shocked into bright star-forming ridges. Tidal forces in NGC2207were perpendicular to its disk,causing perpendicular oscillations and possible warping(Paper II).There is widespread high velocity dispersion in the HI gas in these galaxies,commonly reaching∼50km s−1in otherwise normal-looking regions.These dispersions have since been found in all the ocular interactions(see references above)and in other near collisions(e.g.,Irwin1994).Because of the extraordinarily large HI clouds (108M⊙)that are also present(as observed also in other close encounters–Irwin1994; Smith1994),the observed velocity dispersions promoted the idea that tidal dwarf galaxies can form because of super-size gravitational instabilities in the gas(Elmegreen,Kaufman,&Thomasson1993;Wetzstein&Burkert2005).An even larger cloud located at the tip of the far-western arm in NGC2207and containing108.9M⊙appeared to be triggered by the interaction as well;tidal forces can accumulate outer disk material at this point(Mirabel, Dottori,&Lutz1992;Elmegreen et al.1993;Duc,Bournaud,&Masset2004).NGC2207contains a peculiar region of intense emission at Hαand radio continuum wavelengths(hereafter called Feature i,following the nomenclature in Paper III).Feature i is located on a spiral arm in the western outer disk.Supernova1999ec occurred just south of this region,near two other small IR clumps(see Fig.10below).The Hαluminosity of Feature i,spanning6arcsec in diameter,is equivalent to that of30Dor in the Large Magellanic Cloud.The radio continuum emission comes from a∼1arcsec core with∼300 times the luminosity of Cas A at6-20cm(Vila et al.1990),and an extended component (∼2kpc diameter)with∼1500times the luminosity of Cas A(Paper II).HST reveals an unusual dust cloud elongated for300pc perpendicular to the local spiral arm and containing >106M⊙of gas with>3magnitudes of optical extinction(Paper IV).HST also shows a red V-shaped feature extending for∼500pc in the opposite direction to the dust cloud (Paper III).Between the dust cloud and the red V is a very red region that may contain several massive clusters.Perhaps Feature i is a young star cluster that formed with such a high density that the central stars combined into a black hole,as in models by Ebisuzaki et al.(2001).Spitzer Space Telescope(SST)observations of other galaxies reveal old and young stars at the shortest wavelength(3.6µm)in the Infrared Array Camera(IRAC,see Fazio et al.2004),and PAH emission and heated dust continuum at the longer IRAC wavelengths (4.5µm,5.8µm,and8.0µm;e.g.,Helou et al.2004;Willner et al.2004).The8µm band in IRAC and the longer wavelength bands in the Multiband Imaging Photometer(MIPS;Rieke et al.2004)tend to show moreﬁlamentary structures,such as those seen in M81(Gordon et al.2004a).Here we present IRAC and MIPS observations of IC2163/NGC2207with the SST in order to consider the distribution and properties of clumpy emission at IR wavelengths.The HST data showed the presence of many super star clusters(SSCs),with the most massive having comparable luminosities to the most massive clusters in the Antennae galaxies(Zhang &Fall1999).This is peculiar for IC2163/NGC2207because the total number of clusters is 100times smaller than in the Antennae.Generally,larger samples of clusters have brighter brightest members,with a near-linear relation between brightest luminosity and number of members for an n(L)dL∝L−2dL luminosity(L)distribution function(Whitmore2002). The similarity of the brightest members in IC2163/NGC2207compared with the much more proliﬁc Antennae galaxy suggests that the cluster luminosity function in IC2163/NGC2207is either aﬂatter power law than what is seen in the Antennae,or it is peaked at the high mass end.The most massive clusters are in the eyelid part of IC2163,which,according to our models,is a galactic shock front caused by the sudden radial motion of disk gas responding to tidal torques.Many other clusters have formed inﬁlamentary shocks as well.The SST images are not resolved well enough to separate individual star clusters within a clump.The clumps that can be observed are the equivalent of star complexes,which are collections of OB associations and dense clusters ranging in size from several hundred parsecs to a kpc(Efremov1995).In contrast,our HST observations gave luminosities for individual clusters.We compare these two data sets here.The extended IR emission that was observed with SST will be the subject of a forthcoming paper(Kaufman et al.2005).2.Data and Overall EmissionObservations were made with the Spitzer Infrared Array Camera on February22,2005 and the Multiband Imaging Photometer on March11,2005.For the IRAC observations at3.6,4.5,5.8,and8µm,the full array readout mode was used,with high dynamic range. The frame time was30sec,using a cycling dither pattern with9positions.The images have a scale of1.2arcsec pixel−1,corresponding to∼200pc at35Mpc.For the MIPS observations,the largeﬁeld size was used for24µm and160µm,and the smallﬁeld size for 70µm.For24µm,the exposure time was30sec.For70µm,a raster map was made with a3x3raster and a step size of1/2,with a10sec exposure time.For160µm,4cycles were used,with10sec per cycle.The image scales are2.45arcsec pixel−1,4.0arcsec pixel−1,and 8.0arcsec pixel−1for the24µm,70µm,and160µm bands,respectively.The corresponding point-spread functions are2.0-2.4arcsec(FWHM)for the IRAC images and5.9arcsec,18 arcsec,and40arcsec,respectively,for the MIPS bands.The Basic Calibrated Data images are shown in Figure1,with each band represented separately for IRAC and MIPS.For comparison,a B-band HST WFPC2image is also shown (from Paper IV).The line in the upper left shows a scale of30arcsec,which corresponds to 5100pc at the distance of35Mpc.The clumps and spiral arms in the IRAC bands match the HST optical features well.There is a semi-circular shell1600pc in diameter in the southern arm of NGC2207(marked in Figure2below)that is not obvious in the optical images but partially coincides with an Hαfeature(Fig.3below)and a dust arc.Feature i stands out in all IRAC and MIPS images as an intense source in the western outermost arm of NGC2207.The central2pixels of Feature i are saturated in the24µm image.The eyelid region of IC2163is also prominent in all IRAC and MIPS bands;itdominates the160µm image,although it is unresolved.The MIPS70µm image is similar to the20cm radio continuum image(shown in Paper I).A detailed comparison between MIPS, HI,and radio continuum emissions will be made in a second paper(Kaufman et al.2005).Figure2shows a color composite image of the four IRAC bands,with colors ranging from blue for the short wavelength band to red for the long wavelength band.The prominent features discussed in this paper are identiﬁed.The spiral arms in NGC2207are clearly segmented into regularly spaced beads of emission.The same is true for the oval in IC2163, although the individual beads are diﬃcult to see with the contrast level in this color image (they show up better in Figs.4and15below).Figure3is an overlay of8µm contours on an Hαemission image(from the ground-based continuum-subtracted Hαimage in Paper IV).The bright8µm clumps match the HII regions well.A generally good match between8µm and Hαwas also found for NGC300by Helou et al.(2004).Kennicutt(2005)noted that95%of the infrared clumps in nearby galaxies of the Spitzer Infrared Nearby Galaxy Survey(SINGS–Kennicutt et al.2003)have associated HII regions,and that this implies star-forming regions take less than1Myr to become optically visible.A map made by dividing the3.6µm image by the4.5µm image is shown in Figure4. The nuclei disappear and the star-forming clumps become more prominent.Evidently,the [3.6]/[4.5]ratio is about constant everywhere outside the clumps.This implies that the mixture ofﬁeld stars and diﬀuse hot dust emission is uniform outside the clumps.Inside the clumps,the[3.6]/[4.5]ratio is low(black in the image).This low ratio may be the result of an excess of H Brαat4.5µm from ionization inside the clumps(Churchwell et al. 2004;Whitney et al.2004);the hot and bright stars which contribute to this ionization are not prominent at either3.6µm or4.5µm.Figure5shows scans through both nuclei out to the galaxy edges in3.6µm and4.5µm and in the ratio[3.6]/[4.5].The uniformity of the interclump[3.6]/[4.5]ratio is also evident here.Clumps at the far right of the scan(west in the image)break this uniformity with low values of[3.6]/[4.5].The totalﬂux densities in IRAC and MIPS were determined for the combined galaxies, the eyelid and interior region of IC2163,and Feature i.We used the imstat routine in the Image Reduction and Analysis Facility(IRAF)to measure the totalﬂux inside a box with a ﬁxed size for each region.Backgrounds determined from neighboring boxes were subtracted. The results are given in Table1.For Feature i,the box size was24′′,which is3times the MIPS160µm pixel scale.3.Clumpy Emission3.1.IRAC colorsPhotometry for225IRAC clumps,identiﬁed by eye in the3.6µm image,was done in all possible passbands using phot in IRAF.All the measured clumps are circled in Figure6; stars are denoted by squares.The clumps also show up in the other passbands(Fig1)but the longer wavelengths show the moreﬁlamentary,gaseous structures.The brightest clumps in theﬁgure are numbered and Feature i is indicated;their properties are listed in Table2, including Hαﬂux.The positions and IRACﬂuxes of all the measured clumps are presented in Table3.The photometry was performed using aﬁxed aperture3pixels in radius with background determined from an annulus5pixels in width,separated from the aperture by5pixels.The IRAC counts in MJy/ster were converted toﬂux densities by multiplying by the areas of the pixels,and then converted to magnitudes using the zeropoints given in the online IRAC tables of the Data Handbook1.The photometric accuracy of the IRAC images is better than0.1mag.Weﬁnd that the measurement uncertainties are0.2to0.5mags for diﬀerent regions because of varying sky backgrounds and variations within the clumps.The faintest clump was18.7mag at3.6µm,which corresponds to an absolute magnitude of−14.0at a distance of35Mpc.Color-magnitude diagrams for all measured IRAC clumps are shown on the left in Figure7.Diﬀerent symbols correspond to diﬀerent locations in the galaxies:inner regions (triangles),eyelid or middle(squares),and outer regions(circles).Open symbols are for IC2163andﬁlled symbols are for NGC2207.The rms scatter in color increases as the clumps get fainter.We judged that the color errors start to be signiﬁcant for[3.6]and[4.5] magnitudes fainter than16,[5.8]magnitudes fainter than14,and[8.0]magnitudes fainter plete removal of any object with at least one passband fainter than these limits led to116remaining clumps,which we will refer to as the low-noise group;the other clumps will be referred to as the high-noise group.The color-magnitude diagrams for the 116low-noise clumps are shown on the right of Figure7.The brightest object in all diagrams is Feature i in NGC2207.The next brightest objects are in the eyelid of IC2163,most of which are1-2mag brighter at3.6µm than any other clump.Figures8and9show color-color plots of the116low-noise clumps.They all cluster in a narrow region on these plots,indicating some uniformity to the properties of the interstellarmedium.There are no obvious trends between color and location.In Figure8,the colors of dust(Li&Draine2001)and G0-M5stars(Jones et al.2005)are shown by a plus symbol and diamond,respectively(from Smith et al.2005;Whitney et al.2004).In Figure9,the plus symbol shows the dust again,but the G0-M5stars would be oﬀscale at[4.5]-[5.8]∼0.1 and[5.8]-[8.0]∼0.1.The nuclei of both galaxies are shown by X symbols in bothﬁgures; they have essentially the same colors as G0-M5stars(Jones et al.2005).The[3.6]-[4.5] colors of the clumps are redder than G0-M5stars by∼0.1−0.3mag,signifying populations with average ages younger than A0V stars(e.g.,Pahre et al.2004a).The[4.5]-[5.8]colors are redder than G0-M5stars by∼1.5−3mag,again indicating young ages.Several studies have considered whether IR emission from a given region is starlight-dominated,PAH-dominated,or blackbody emission in certain passbands.Engelbracht et al.(2005)examined IR spectra,IRAC,and MIPS data of spiral galaxies,compared them with stellar evolution models,and showed a clear division between sources containing PAH emission and sources not containing PAH emission based on the ratio of starlight subtracted 4.5µmﬂux density to8µmﬂux density;a ratio less than0.05corresponds to the presence of PAH emission.They assumed that the3.6µm emission is dominated by old starlight(e.g., Helou2004).Following their method,we subtracted a scaled3.6µm emission(with a factor 0.57which they based on Starburst99models from Leitherer et al.1999)from the4.5µm emission for each clump;the ratio of this reduced4.5µm emission to the8µm emission fell below0.05for215out of the225clumps,indicating that these clumps should contain PAH emission.Such PAH emission is expected to contribute mostly to5.8µm and8µm bands (e.g.Helou et al.2004,Wang et al.2004).Sajina,Lacy,&Scott(2005)showed that PAH’s are located in a diﬀerent region of log[S8/S4.5]versus log[S5.8/S3.6]than starlight(see their Fig.6).Our clumps are almost all in the range of log(S8/S4.5)=0to1,log(S5.8/S3.6)=0.7 to1.5for theseﬂux ratios,which falls within the PAH regime of the diagram.Also,our clump color distribution for[5.8]-[8]peaks between1.6and2.2,which is consistent with PAH-dominated emission at these wavelengths,as modeled by Li&Draine(2001).In other SST observations of interacting galaxies,Arp107,at a distance of138Mpc,was found to have numerous clumps in a ring-like distribution with a1.3mag azimuthal variation of[4.5]-[5.8]colors indicating a stellar age gradient,as matched by numerical simulations (Smith et al.2005).The clumps studied in Arp107were divided into brighter and fainter ones.The brighter clumps had slightly redder[4.5]-[5.8]and[5.8]-[8.0]colors(by about0.5 mag)than the fainter clumps.Smith et al.interpret the redder clumps as having a higher proportion of young to old stars.Our clumps on average have about the same[3.6]-[4.5] colors as all of their clumps,∼-0.2to0.3,with colors that are about0.5mag redder than their bright clumps in[4.5]-[5.8]and similar to their bright clumps in[5.8]-[8.0].These redder colors suggest that our regions are slightly younger on average than theirs.Spitzer IRAC observations of the Antennae galaxy pair NGC4038/4039,at a distance of21Mpc,reveal clumpy star formation and extensive PAH emission(Wang et al.2004). Their average[5.8]-[8.0]clump color is∼1.8mag,which is consistent with our clump colors. M51also has prominent star-forming clumps and brightﬁlaments in SST images(Calzetti et al.2005).IRAC and MIPS observations of M81(Gordon et al.2004a)and M31(Gordon et al.2004b)show that the prominent clumpy UV,8µm,and24µm features correspond well to one another.The alignment between IRAS clumps and Hαemission in IC2163/NGC2207,the color-color distribution of the IRAC clumps,and the likely presence of PAH emission from them,all suggest that these regions are relatively young.Precise age estimates are not possible without better calibrations of the relative contributions to the IRAC bands from young and old stars, dust,free-free emission,PAHs,and other possible emission sources.Nevertheless,the young ages found here are consistent with the ages derived by more conventional techniques using HST data,which are up to several×107years(Paper IV).3.2.Sample Comparisons with HST emissionEach IRAC clump is composed of numerous smaller clusters and associations that are individually visible with HST.The super star clusters found with HST,for example,are inside the IRAC clumps,although these IRAC clumps do not appear to be more luminous or special in any way.Figure10shows two IRAC star forming regions in detail;Feature i is on the left and clump IR9from the eyelid region is on the right.The contours are8µm emission at2.2 arcsec resolution and the gray scale is the optical emission at0.13arcsec resolution.In the vicinity of Feature i,there are two∼104.5M⊙super star clusters(SSC)from the catalog of Paper IV,many other clusters and associations,four IRAC sources from Table3,an optical supernova source,and a dust streak to the southwest of a very red cluster near the center of the large8µm source.In the vicinity of IR9,there are2super star clusters and2other IRAC sources.parison with HαemissionThe luminosities of HII regions were calculated in Paper IV based on ground-based Hαimages.The Hαresolution is similar to the IRAC resolution.For individual clumps there is a correlation between the Hαluminosity,which ranges from5×104L⊙to4×106L⊙,and the passband-combined IRAC luminosity,which ranges from2.5×106L⊙to2.5×108 L⊙(using L⊙=4×1033erg s−1).There is a scatter around this correlation of∼0.5mag, which is most likely the result of slight age and extinction diﬀerences.The luminosities are compared in Figure11.The brightest feature in both passbands is again Feature i.For the individual clumps,the average ratio of Hαto3.6µm is0.052±0.062.For Arp107,this ratio averages0.04(Smith et al.2005).For Hα/8µm,our average ratio is0.005±0.006and for Arp107,the ratio ranges from0.006to0.05.For the combined galaxies,the ratio of Hαto3.6µm luminosity is0.009.This ratio is slightly less than that in Arp107,where it is0.012for the large galaxy and0.025for the small galaxy(Smith et al.2005).It is also slightly less than in NGC7331,where it is0.013 (Regan et al.2004).The ratio of Hαto8µm for the combined IC2163/NGC2207pair is 0.007,which is slightly less than0.0083and0.025for Arp107,and0.030for NGC7331.3.4.Spectral energy distributionsThe IRAC SEDs for the clumps in Table2are shown in Figure12,with the log of the wavelength timesﬂux density,λFλ,versus the log of the wavelength.The B,V,and Iﬂux densities from HST are also plotted.They were measured in the HST images using apertures 36pixels in radius to correspond to the apertures used for the IRAC images.Some of the SEDs dip down toward the blue,presumably because of extinction in these optical bands (see Paper IV).All of the clumps show a dip at4.5µm.The brightest clump is Feature i,which is represented twice in Figure12:once as an IRACﬂux like the other clumps and again as aﬂux from all passbands inside a24′′box, including MIPS images.The line with open circles indicates the IRAC values using the usual 1.2′′pixel size of this instrument,with photometry done in an aperture of3pixel radius,as for the other clumps.Theﬂux inside the24′′box is indicated by a dotted line and x-marks; it has the background subtracted as determined from neighboring boxes.The two curves for Feature i are similar at short wavelengths because the3-pixel central source dominates the emission inside the24′′box.The combined light from IC2163and NGC2207is shown in Figure12as the long solid line.The eyelid plus interior of IC2163is shown as a long dashed line.The integrated SEDs from other galaxies observed in the SINGS survey are similar to the integrated SED in Figure12.Dale et al.(2005)show SEDs for galaxies spanning a wide range of Hubble types.Our combined-galaxy SED resembles the integrated SEDs for the Seyfert I galaxy NGC1566(SABbc),the HII galaxy NGC2403(SABcd),and the LINERgalaxies NGC1097(SBb),NGC7552,and NGC3190(SAap).These galaxies all have a dip at4.5µm,a local peak at5.8or8µm,a trough at24µm,and a rise to70µm or160µm,like IC2163/NGC2207combined.There are also other galaxies in Dale et al.(2005)that are qualitatively similar;many of them are LINERs or Seyferts.Theoretical curves in Dale et al.(2005)ﬁt the region from24µm to160µm with con-tinuum emission from dust.We determined dust temperatures from MIPS data for the combined galaxies,the eyelid and interior regions of IC2163,and Feature i.They are based on theﬂux densities in Table1.Because there can be diﬀerent dust components contribut-ing to the short and long wavelength bands of MIPS,we calculated blackbody temperatures from both the24µm/70µm ratio and the70µm/160µm ratio.Dust temperatures were not determined for individual clumps because of the lack of resolution in the long wavelength bands.For these three regions,the dust temperatures from the short wavelength ratios were approximately50K,52K,and58K,respectively.Similarly,the dust temperatures from the long wavelengthﬂux ratios were25K,24K,and36K,respectively.The uncertainties in ourﬂux densities are about20%based on the MIPS Data Handbook2.This would give a temperature range of±3K for the above measurements.Thus,the temperatures for the combined galaxies and the eyelid are essentially the same,but the temperature in Feature i is slightly warmer.This is also evident from the SED in Figure12,which shows a downturn from70µm to160µm for Feature i,unlike the other SEDs.3.5.IRAC Luminosity FunctionsClump luminosities were determined from the IRAC bands.The MIPS bands contain diﬀuse ISM emission that is unrelated to star complexes(Haas et al.2002),and they are also poorer resolution,so we do not include them here.For the IRAC bands we integrated theﬂux density over wavelength,L IRAC erg s−1 =4πD210−234 i=1F i∆λi,(1) where the bandwidths∆λfor the4bands are0.75µm,1.015µm,1.425µm,and2.905µm3; i=1corresponds to3.6µm,i=2corresponds to4.5µm,etc.,and F i is theﬂux density in Jy per micron(=10−4Fνc/λ2with cgs units for speed of light c and wavelengthλ).For this equation,D is the distance of35Mpc,evaluated in cm.Figure13shows the IRAC luminosity functions for the three main regions of each galaxy (inner,middle,and outer,as in Figs.7-9).The bar-chart histograms are counts in bins of width0.5in the log of the luminosity(in L⊙)for the low-noise clumps.The open squares are counts for all the clumps,including the high-noise clumps.The x-marks on the abscissae are the sensitivity limits,106.57L⊙,deﬁned by an integration over the separate limits mentioned in reference to Figure7:16th mag for[3.6]and[4.5],14th mag for[5.8],and12.5mag for [8].The smooth curves in Figure13calculate the luminosity distribution in a diﬀerent way, this time including both the low-noise group(116clumps,plotted with solid-line curves)and the combined low and high-noise groups(116+94clumps,plotted with dashed curves).In this method,we add a Gaussian distribution for each clump with an area equal to unity and a dispersion equal to the uncertainty in the luminosity obtained from the uncertainty in the photometry measurement(as calculated by phot in IRAF).This error is low for the brighter clumps(0.16of a magnitude at log L⊙=7.9)and relatively high for the faintest clumps(0.3of a magnitude at log L⊙=6.57,which is the limiting sensitivity).In this way, the fainter clumps can be included with their high uncertainties.The resultant summed distributions were then multiplied by4to give them a height in Figure13comparable to the histogram.The Gaussian-sum distributions resemble the binned distributions for the low-noise clumps(the dotted histograms),and extend to lower luminosity for the high-noise clumps,resembling the counts shown by open squares.The dots in the Gaussian-sum curves represent the centers of the bins used for counting.The combined luminosity functions for all of the IRAC sources are presented in Figure 14.The panel on the left is for the sum of the IRAC bands,and the panel on the right is for the3.6µm band only,calculated as∆λ3.6F3.6and converted to cgs.As in the previousﬁgure, the solid curves are for the low-noise sources and the dashed curves are for all the sources including those with high-noise.Dots on the curves denote the centers of the counting bins. Feature i is the single peak on the right in the left panel;it appears only in the combined luminosity function.The open squares are counts for all the sources in intervals of0.5mag. The two sensitivity limits are shown as x-marks on the abscissae.The combined IRAC luminosity function(Fig.14)is steeper than the usual power law with a slope of−1on such a log-log plot.In Paper IV we also found that the HII region luminosity function has a steep drop-oﬀat the upper end.A high fraction of clumps are in thinﬁlaments,including the whole eyelid region and many of the arcs of spiral arms.In these regions,the complexes look like regular beads on a string(e.g.,Figs.2and4).Figure15illustrates several of the regions where the clumps appear evenly spaced.The8µm logarithmic image is shown because this band has a。

翻译——精选推荐

细胞分子生物学文章第十卷（2005），711-719 pl2005.7.15寄出200510.6收到脂质体：一项先进制造技术的概述新西兰，北帕默斯顿，专用邮袋11222，梅西大学Riddet中心，M.REZA MOZAFARI摘要：近几十年来，脂质体作为生物膜的理想模型，也是药物、诊断、疫苗、营养物和其他生物活性剂的有效载体，引起了广泛关注。

在不同背景下研究者们对脂质体学领域的文献报道广泛地不断地增加，这表明这一领域引人入胜。

自从大约40年前脂质体被介绍到科学界，许多技术和方法在或大或小的脂质体制造规模上得到发展。

这篇文章将在大体上提供脂质体制备方法优缺点的概览，特别强调在我们实验室开发的加热法，作为一种脂质囊泡快速生产的模式技术。

关键词：载体系统，加热法，脂质囊泡，脂质体学，制造技术引言脂质体科学技术是一个正在飞速发展的科学，举几个例子，它用于诸如药物递送，化妆品，生物膜的结构和功能，探索生命起源等领域。

这是由于脂质体有一些有利的特性，例如，它不仅能包含水溶性药物也能包含脂溶性药物，在体内识别特定靶向位点，在流动性、大小、电荷、层数的方面具有多样性。

脂质体作为生物膜模型的应用限于在实验室中研究，它们在生物活性剂的包载和递送的成功应用不仅取决于脂质体载体可以达到预期目的的优越性的示范，还取决于技术和经济可行性的规划。

对于递送应用，脂质体配方应该具有高包封率，窄粒度分布，持久稳定性和理想的释放特性（根据预期的应用）。

这些要求制备方法有产生脂质体的可能性，且脂质体可采用多种成分分子，例如：脂质/磷脂可提高脂质体稳定性。

除了上述特性，对于蛋白质、核酸之类敏感的分子/化合物的递送，脂质体也应该能保护复合制剂，防止其退化。

尽管在脂质体上进行了大量的研究开发工作，但只有少数脂质体产品已被批准为人类使用至今。

这也许有许多原因：一些脂质体配方的毒性，分子和化合物在脂质体中的低包封，脂质体载体的不稳定性，脂质载体的不稳定性，特别是大尺度的脂质体生产成本高。

X-ray Emission from Haloes of Simulated Disc Galaxies

a r X i v :a s t r o -p h /0201529v 1 31 J a n 2002Mon.Not.R.Astron.Soc.000,1–8(2001)Printed 1February 2008(MN L A T E X style ﬁle v2.2)X-ray Emission from Haloes of Simulated Disc GalaxiesS.Toft,1⋆,J.Rasmussen 1,J.Sommer-Larsen 2and K.Pedersen,11Astronomical Observatory,Copenhagen University,Juliane Maries Vej 30,DK-2100Copenhagen Ø,Denmark2TheoreticalAstrophysics Center,Juliane Mariesvej 30,DK-2100Copenhagen Ø,DenmarkABSTRACTBolometric and 0.2-2keV X-ray luminosities of the hot gas haloes of simulated disc galaxies have been calculated at redshift z =0.The TreeSPH simulations are fully cosmological and the sample of 44disc galaxies span a range in characteristic circular speeds of V c =130-325km s −1.The galaxies have been obtained in simulations with a considerable range of physical parameters,varying the baryonic fraction,the gas metallicity,the meta-galactic UV ﬁeld,the cosmology,the dark matter type,and also the numerical resolution.The models are found to be in agreement with the (few)relevant X-ray observations available at present.The amount of hot gas in the haloes is also consistent with constraints from pulsar dispersion measures in the Milky Way.Forthcoming XMM and Chandra observations should enable much more stringent tests and provide constraints on the physical parameters.We ﬁnd that simple cooling ﬂow models over-predict X-ray luminosities by up to two orders of magnitude for high (but still realistic)cooling eﬃciencies relative to the models presented here.Our results display a clear trend that increasing cooling eﬃciency leads to decreasing X-ray luminosities at z =0.The reason is found to be that increased cooling eﬃciency leads to a decreased fraction of hot gas relative to total baryonic mass inside of the virial radius at present.At gas metal abundances of a third solar this hot gas fraction becomes as low as just a few percent.We also ﬁnd that most of the X-ray emission comes from the inner parts (r <∼20kpc)of the hot galactic haloes.Finally,we ﬁnd for realistic choices of the physical parameters that disc galaxy haloes possibly were more than one order of magnitude brighter in soft X-ray emission at z ∼1,than at present.Key words:methods:N-body simulations –cooling ﬂows –galaxies:evolution –galaxies:formation –galaxies:halos –galaxies:spiral –X-rays:galaxies1INTRODUCTIONIn disc galaxy formation models infall of halo gas onto the disc due to cooling is a generic feature.However,the gas accretion rate and hot gas cooling history are at best uncer-tain in all models so far.It is thus not clear to which extent the gas cooling out from the galaxy’s halo is replenishing that which is consumed by star formation in the disc.Such continuous gas infall is essential to explain the extended star formation histories of isolated spiral galaxies like the Milky-Way and the most likely explanation of the “G-dwarf prob-lem”—see,e.g.,Rocha-Pinto &Marciel (1996)and Pagel (1997).At the virial temperatures of disc galaxy haloes the dominant cooling mechanism is thermal bremsstrahlung plus atomic line emission.The emissivity,increasing strongly with halo gas density,is expected to peak fairly close to the disc and decrease outwards,and if the cooling rate is signif-⋆E-mail:toft@astro.ku.dkicant the X-ray ﬂux may be visible well beyond the optical radius of a galaxy.Recently,Benson et al.(2000)compared ROSAT ob-servations of three massive,nearby and highly inclined disc galaxies with predictions of simple cooling ﬂow models of galaxy formation and evolution.They showed that these models predict about an order of magnitude more X-ray emission from the galaxy haloes (speciﬁcally from a 5-18arcmin annulus around the galaxies)than observational de-tections and upper limits.In this paper we present global X-ray properties of the haloes of a large,novel sample of model disc galaxies at red-shift z =0.The galaxies result from physically realistic grav-ity/hydro simulations of disc galaxy formation and evolution in a cosmological context.We ﬁnd that our model predic-tions of X-ray properties of disc galaxy haloes are consistent with observational detections and upper limits.Given the results of the theoretical models of Benson et al.we list the most important reasons why simple cooling ﬂow modelsc2001RAS2S.Toft et al.-60-40-200204060x [kpc]34567l o g T [K]Figure 1.The ﬁgure shows the temperature of the SPH gas particles in a typical disc galaxy from our simulations versus their x coordinate (one of the axes in the disc).The “cold”(log T <4.5)gas which is primarily situated in the disc is removed from the catalogues since it does not contribute to the X-ray ﬂux.over-predict the present day X-ray emission of disc galaxy haloes.In section 2we give a very short description of the disc galaxy simulations.In section 3we brieﬂy describe the X-ray halo emission calculations and in section 4the results obtained.Section 5constitutes the discussion and section 6the conclusion.2DISC GALAXY SIMULATIONSWe have in recent years developed novel models of formation and evolution of disc galaxies.The model disc galaxies result from ab initio ,fully cosmological (ΩM =1or ΩM +ΩΛ=1),gravity/hydro simulations.These simulations are started at a suﬃciently high redshift (z i =20-40)that the density per-turbations are still linear and are then evolved through the entire non-linear galaxy formation regime to the present epoch (z =0).The code uses a gridless,fully Lagrangian,3-D TreeSPH code incorporating the eﬀects of radiative cooling and heating (including the eﬀects of a meta-galactic UV ﬁeld),inverse Compton cooling,star formation,and ener-getic stellar feedback processes.A major obstacle in forming realistically sized disc galaxies in such simulations is the so-called “angular mo-mentum problem”(e.g.,Navarro &White 1994,Sommer-Larsen et al.1999).We overcome this problem in two dif-ferent ways:a)Using cold dark matter (CDM)+stellar feedback processes (Sommer-Larsen et al.2002)or b)Us-ing warm dark matter (WDM)(Sommer-Larsen &Dolgov 2001).A total of 44such disc galaxy models with charac-teristic circular velocities in the range V c =130–325km s −1form the basis of the predictions presented in this paper.The simulations initially consist of 30000-400000SPH+DM particles and in the majority of them some of the SPH par-ticles are turned into star particles over the course of theFigure 2.Bolometric luminosity as a function of characteristic circular speed.Small symbols :Flat ΩM =1.0cosmology:Open symbols:baryon fraction f b =0.05,ﬁlled circles f b =0.1.Triangles:without UV ﬁeld,non-triangles:with a UV ﬁeld of the Efstathiou (1992)type.Connected symbols are the same galaxies run with medium (open circles)and high (open circles with crosses)res-olution.All simulations represented by small symbols have pri-mordial rge symbols :Flat (ΩΛ,ΩM )=(0.7,0.3)cosmology:Open symbols:f b =0.05,ﬁlled symbols f b =0.1.Cir-cles correspond to primordial abundance and with a Haardt &Madau (1996)UV ﬁeld,squares correspond to Z =1/3Z ⊙(us-ing the cooling function of Sutherland &Dopita 1993,which does not include eﬀects of a UV ﬁeld).The curves are the L X,bol -V c relationship for the simple cooling ﬂow models (for ΛCDM NFW haloes)described in Sec.5.The curves represent diﬀerent bary-onic fractions (solid curves have f b =0.1,dotted curves have f b =0.05)and abundances (thick curves:primordial abundances,thin curves:Z =1/3Z ⊙).simulation —for details about the TreeSPH simulations we refer the reader to the above quoted references.Most of the simulations were run with primordial gas composition (76%H and 24%He by mass)under the as-sumption that the inﬂowing,hot gas is fairly unenriched in heavy elements.To test the eﬀects of metal abundance eight ΛCDM simulations (four with “universal”baryonic fraction f b =0.05and four with f b =0.10)were run with a gas abundance of 1/3solar (speciﬁcally [Fe/H]=-0.5).This is the metal abundance of the intracluster medium and can probably be considered a reasonable upper limit to the metal abundance of the hot gas in disc galaxy haloes.3X-RAY HALO EMISSION CALCULATIONFor each of the simulated galaxies at z =0we create a cata-logue of SPH gas particle positions,densities,temperatures and masses.For each catalogue a box of size (1000kpc)3centered on the galaxy is retained and all gas particles with temperatures log(T)<4.5are cut away (see Fig.1)sincec2001RAS,MNRAS 000,1–8X-ray Emission from Haloes of Simulated Disc Galaxies3Figure 3.0.2-2keV band luminosity as a function of character-istic circular speed.Symbols as in Fig.2.they will not contribute to the X-ray ﬂux (these are mainly gas particles which have cooled onto the disc).The density and temperature of each particle is then averaged over itself and its ﬁve nearest neighbors using a spherical smoothing kernel proposed by Monaghan &Lattanzio (1985)1,and a volume is assigned to the particle given its mass and ing the average temperatures and densities,each SPH particle is treated as an optically thin thermal plasma,and the associated X-ray luminosity is calculated at the rele-vant position in a given photon energy band with the meka plasma emissivity code (Mewe et al.1986).X-ray luminosi-ties are then computed by summation over all particles in the volume of interest.The bolometric X-ray luminosity is calculated using the 0.012−12.4keV band.4RESULTSIn Fig.2the total bolometric X-ray luminosities L X,bol of the 44simulated disc galaxies in our sample are plotted versus their characteristic circular speed V c ,deﬁned as the circu-lar velocity in the disc at R 2.2=2.2R d ,where R d is the disc scale length –see Sommer-Larsen &Dolgov (2001)for details.The X-ray luminosities derived from the simulations are up to two orders of magnitude below values derived from simple cooling ﬂow models which are described in Sec.5.1.As can be seen from the ﬁgure L X,bol ∼1040erg s −1for a Milky Way sized galaxy.As expected from the simple models,the simulated1It is important to note that we average over the original den-sities from the TreeSPH simulations.One can show that if the cold gas particles are cut away and the densities of the remaining gas particles are then recalculated using the full SPH procedure the X-ray luminosities will be underestimated due to resolution problems at the disc–halo interface.1020304050|z| [kpc]10-1810-1710-1610-1510-1410-1310-12S (z ) [ e r g s -1 a r c m i n -2 c m -2]305 < V C < 325202 < V C < 240Figure 4.Mean bolometric X-ray surface brightness proﬁles per-pendicular to the plane of the disc of three high V c galaxies (solid curve)and four Milky Way sized galaxies (dashed curve).The physically most important parameters for these galaxies are f b =0.10and primordial gas abundance.The proﬁles were calcu-lated by binning the emission in 40kpc wide,5kpc high slices parallel to the disc.galaxies display an L X,bol -V c relation,but with a signiﬁcant scatter.Part of this scatter arises from the diﬀerent condi-tions under which the simulations have been run (see Sec.5.3)and part of it is a “real”scatter arising from the diﬀer-ent geometries and cooling histories of the individual galaxy haloes.This “real”scatter can be estimated by inspection of the large ﬁlled circles in Fig.2which represent simulations run with the same,physically important parameters (see the ﬁgure caption).These data points display a rms dispersion of about 50%around the mean.There is a tendency for galaxies formed in simulations with baryon fraction f b =0.05and primordial gas abundance to have systematically higher L X,bol (by about a factor of two)than galaxies formed in similar simulations with f b =0.10.Also,galaxies formed in simulations with gas abun-dance Z =1/3Z ⊙tend to have systematically lower L X,bol than the ones with primordial gas.We discuss these trends in Sec.5.Fig.3shows the 0.2-2keV band X-ray luminosities of the 44sample disc galaxies versus V c .The systematic trends mentioned above are also seen in this plot,in particular is the diﬀerence between the f b =0.05and 0.10simulations (with primordial gas)even more pronounced than in Fig.2.We also discuss this in Sec.5.Most (but not all)of the X-ray emission originates from regions of the hot gas halo fairly close to the disc:95%of the emission typically originates within about 20kpc of the disc.This is illustrated in Fig.4where we plot the mean surface brightness proﬁles perpendicular to the disc of three high V c galaxies and four Milky Way sized galaxies.5DISCUSSIONc2001RAS,MNRAS 000,1–84S.Toft et al.5.1Comparison to simple coolingﬂow modelsIn order to compare our results with previous work,we cal-culated a family of simple coolingﬂow models similar to those considered by Benson et al.(2000).In these models it is assumed that the cooling occurs in a static potential and that the gas initially was in place and traced the dark matter(DM).Gas is assumed toﬂow from the cooling radius(at which the cooling time equals the age of the universe)to the disc(settling there as cold gas)on a time-scale much shorter than the Hubble time.The bolo-metric X-ray luminosity can then be approximated simply as the mass accretion rate,˙M cool=4πr2coolρgas(r cool)˙r cool, times the gravitational potential diﬀerence,soL X,bol=˙M cool(r cool) r cool r optical V2c(r)X-ray Emission from Haloes of Simulated Disc Galaxies5L X=45±4.4×1040erg s−1for NFW haloes(and even more for isothermal sphere haloes).However,the above result is in excellent agreement with the expectation from our simu-lations,as illustrated in Fig.5where we plot the bolometric luminosity in the considered annulus(assuming a distance of13.8Mpc)of all the galaxies in our simulation sample ver-sus V c.We note that our three“data”points at V c=300-325 km s−1are for simulations with baryon fraction f b=0.10and primordial gas abundance.For a given V c we would expect the bolometric luminosities to be about twice as much for simulations with f b=0.05-see section5.3.4.Such a low f b is unlikely given determinations of the“universal”baryon fraction,f b≈0.10(h/0.7)−3/2,derived from galaxy clusters (Ettori&Fabian1999)and galaxy clustering(Percival et al. 2001),but in any case the X-ray luminosities would still be consistent with the NGC2841measurement.Moreover,the (more realistic)inclusion of some level of enrichment in the gas and the use of the more realistic meta-galactic UVﬁeld of Haardt&Madau(1996)would tend to lower the X-ray lu-minosities.We also note that the total X-ray luminosities of our three galaxies at V c=300-325km s−1are“only”a factor of3-5lower than the predictions of the simple coolingﬂow models of Benson et al.and ours(the latter for f b=0.10and primordial gas).Hence an important part of the reason why our models match the observed bolometric luminosity of the 5-18arcmin annulus around NGC2841is that most of the X-rays are emitted from the inner20kpc of our model disc galaxy haloes(5arcmin correspond to20kpc at a distance of13.8Mpc).In other words our“geometric correction”is considerably larger than the one used by Benson et al.The observational constraints on the bolometric lumi-nosity of NGC4594and NGC5529are weaker than for NGC 2841.The upper limits on these are again about an order of magnitude less than expected from the simple models,but in agreement with the expectation from our simulations.The diﬀuse X-ray luminosity of the Milky Way’s hot halo has recently been estimated:Pietz et al.(1998)es-timate a0.1-2keV luminosity of7·1039erg s−1and Wang (1997)a0.5-2keV luminosity of3·1039erg s−1.Assuming a temperature of0.15keV(Georgantopoulus et al.1996;Par-mar et al.1999)this translates into0.2-2keV luminosities of5and7·1039erg s−1respectively.These estimates(which should probably be seen as upper limits)are consistent with ourﬁndings from the simulations for V c≃220km s−1—see Fig.3.5.3Eﬀects of numerical resolution and physicalparametersThe galaxies in our sample have been compiled from a number of simulations which have been run with diﬀer-ent cosmological and environmental parameters.These in-homogeneities in our simulation sample may introduce some scatter in the L X−V c diagram.On the other hand,this allows us to investigate trends when varying the physical parameters.In general,varying a parameter has impact on the present day X-ray luminosity of a given galaxy if it signif-icantly alters the ability of the hot halo gas to cool during its life time.In the simple coolingﬂow models,increasing the cooling eﬃciency leads to an increase in L X,while this is not necessarily the case in more realistic simulations.If the cooling eﬃciency is increased,more gas has cooled out on the disc at z=0.This results in an increase of the char-acteristic circular speed of the disc as its dynamics become more baryon(and less DM)dominated,and usually a de-crease in the X-ray luminosity of the halo since there is less hot gas left in the halo to cool and contribute to the X-ray emission.In the following we brieﬂy discuss how the derived re-sults depend on the resolution of the simulation,the pres-ence of an external UVﬁeld,the assumed gas metallicity, the baryon fraction f b,the cosmology,the dark matter type and whether or not star formation is incorporated in the simulations.5.3.1Eﬀects of resolutionAn important test of all numerical simulations is to check whether the results depend on the resolution.This can be done from Fig.2by comparing the connected symbols. These represent the same galaxy,run with normal(open symbols)and8times higher mass+2times higher force resolution(open symbols with crosses).It can be seen that this signiﬁcant increase in resolution only leads to a very modest increase of19±64%in the X-ray luminosity rela-tive to the mean L X−V c relation(i.e.taking the eﬀect of the change of V c with resolution into account).A similar result is inferred from Fig.3.5.3.2Eﬀects of a meta-galactic UVﬁeldThe main eﬀects of a hard UV photonﬁeld is to ionize the gas,signiﬁcantly reducing its ability to cool by colli-sional excitation(line-cooling)mechanisms(Vedel,Hellsten and Sommer-Larsen1994).This eﬀect is evident in Fig.2 where the set of4small triangles and the set of4small open circles represent the same4haloes,run under exactly the same conditions,except that for the latter eﬀects of a meta-galactic,redshift-dependent UVﬁeld were included in the cooling/heating function.There is a tendency for the galaxies without an external UV-ﬁeld to have lower L X,bol and higher V c than the galaxies with external UV-ﬁeld:The former have a median bolometric luminosity of55±16%of the latter(again taking into account the change of V c).So in this case an increase in the cooling rate leads to a decrease in L X,bol at the present epoch since a smaller amount of hot gas is left in the halo to produce the emission –see Fig.6.Note however that in the above simulations with UVﬁeld we used aﬁeld of the Efstathiou(1992)type which is too hard and intense compared to the more realistic one of Haardt&Madau(1996)—see also Sec.5.3.5.Hence,the suppression in cooling eﬃciency and the related increase in L X,bol for disc galaxy haloes formed in simulations with a Efstathiou type UVﬁeld is somewhat too large.5.3.3Eﬀects of gas metallicityThe eﬀects of the metallicity of the gas on the derived L X,bol can be investigated by comparing the squares in Fig.2(which have Z=1/3Z⊙)with the rest of the symbols(which have primordial abundance).In the simple coolingﬂow models, an increase of the metal abundance leads to an increase inc 2001RAS,MNRAS000,1–86S.Toft etal.Figure 6.Bolometric luminosity relative to the one expected for the simple cooling ﬂow models versus hot gas fraction -see text for details.Symbols as in Fig.2.the cooling rate and L X,bol (compare the thick and thin curves in Fig.2);however this is not what we ﬁnd from our simulations.The galaxies with Z =1/3Z ⊙have systemat-ically lower L X,bol than the galaxies with primordial abun-dance,by about a factor of 3-4for the same f b .This is in agreement with the argument that increasing the cooling ef-ﬁciency leads to a decrease in L X,bol at z =0.Note that for the Z =1/3Z ⊙simulations we used the cooling function of Sutherland &Dopita (1993)which does not include the eﬀect of a UV ﬁeld.So we can not completely disentangle the eﬀects of gas metal abundance versus lack of UV ﬁeld on the X-ray luminosities,but Fig.6strongly hints that the former is the most important.5.3.4Eﬀects of baryon fractionComparing in Fig.2L X,bol for simulations run with primor-dial gas and with f b =0.05and 0.10,respectively,the former are more luminous on average by about a factor of two.Yet again we see how increased cooling eﬃciency leads to a de-creased L X,bol at z =0.Summarizing the above results we ﬁnd no statistically sig-niﬁcant dependence of the X-ray luminosities of the simu-lated galaxies on numerical resolution.With respect to cool-ing eﬃciency there is a general trend of higher cooling ef-ﬁciency over the course of the simulation to result in less hot gas left in the halo at z =0to cool,yielding a lower X-ray luminosity.This is qualitatively demonstrated in Fig.6,which shows the bolometric luminosity relative to the one expected for the simple cooling ﬂow models versus the hot gas fraction f (hot gas )=M (hot gas )(<r vir )/(f b M vir ),where M (hot gas )(<r vir )is the mass of hot gas (log(T )>4.5)insideof the virial radius r vir and M vir is the total mass (baryonic +DM)inside r vir .It is seen that only for hot gas frac-tions f (hot gas )>∼0.4-0.5,requiring a physically implausible parameter combination of f b =0.05,primordial gas and an unrealistically hard and intense UV ﬁeld,can our models match the L X,bol predicted by the simple cooling ﬂow mod-els.Pulsar dispersion measures can be used to place ob-servational upper limits on the amount of hot gas in the halo of the Milky Way.We ﬁnd from our simulations that Milky Way sized galaxies formed in primordial gas simu-lations have about 109M ⊙of hot gas inside of 50kpc at z =0and the ones from the Z =1/3Z ⊙simulations about 108M ⊙.Both values are consistent with the observational upper limits of about 2·109M ⊙from pulsar dispersion mea-sures to the Magellanic Clouds and the globular cluster M53(Moore &Davis 1994;Rasmussen 2000).The diﬀerence between the f b =0.05and 0.10cases is even more pronounced for the 0.2-2keV X-ray luminosities,as shown in Fig.3.The reason is that at a given charac-teristic circular speed V c the dynamics of the inner galaxy (where V c is determined)are more baryon dominated for f b =0.10than for f b =0.05.This in turn means that the hot halo is smaller and cooler for the f b =0.10case than for the f b =0.05case.This is demonstrated in Fig.7,which shows the average temperature of the central hot halo gas (inside of 20h −1kpc)for the 44simulations.At a given V c the tem-perature of the inner,hot halo is systematically shifted to lower values for f b =0.10as compared to f b =0.05.Hence for the relevant,relatively low temperatures (T <∼0.3keV,orequivalently,V c <∼300km s −1)less of the emitted radia-tion has energies above 0.2keV for the former than for the latter case.On the issue of baryon dominance of the inner galaxy dynamics note that for a given DM halo,f b =0.10(as compared to f b =0.05)leads to a larger V c and a smaller d v c (R )/d R in the outer parts of the disc,where v c (R )is the rotation curve –this is in line with the ﬁndings of Persic,Salucci &Stel (1996)on the basis of a large observational sample of disc galaxy rotation curves.Finally,as mentioned in Sec.5.2,the temperature of the Milky Way’s inner halo is about 0.15keV corresponding to 1.5-2·106K in agreement with Fig.7.5.3.5Eﬀects of dark matter type,cosmology and star formationWe do not ﬁnd any dependence on the DM type,i.e.whether the simulations are of the WDM or CDM +feedback type.Neither do we ﬁnd any indications of systematic trends with cosmology (SCDM/WDM versus ΛCDM/WDM)although more overlap in ﬁgures 2and 3between the two cosmologies would have been desirable (the reason for this lack of overlap is that rather diﬀerent cosmological volumes were sampled in the two cosmologies —the box size was ∼40h −1Mpc for the ΩM =1cosmology and 10h −1Mpc for the Λ-cosmology).Note also that in Fig.6the results for the diﬀerent cosmolo-gies fall along the same continuous sequence.The reason why the disc galaxies formed in primordial gas simulations for the Ωm =1cosmology tend to have slightly higher relative bolometric luminosity and f (hot gas )than the Λ-cosmology ones is most likely due to the two diﬀerent models of thec2001RAS,MNRAS 000,1–8X-ray Emission from Haloes of Simulated Disc Galaxies7Figure 7.Average inner halo hot gas temperature versus char-acteristic circular speed -see text for details.Symbols as in Fig.2.meta-galactic UV ﬁeld used:For the former we used the one suggested by Efstathiou (1992)(see Vedel et al.1994),whereas for the latter we used the more realistic one from Haardt &Madau (1996).The former has z reionization =∞and is considerably harder and,at z >∼2,more intense than the latter which has z reionization ≃6.Finally,we do not ﬁnd any dependence on whether star formation is included or not in the simulations.5.4Mass accretion ratesIn Sommer-Larsen et al.(2002)disc gas accretion rates due to cooling-out of hot gas are determined for the ΛCDM +feedback model disc galaxies considered there (a sub-set of the sample considered here obtained in simulations with f b =0.10,primordial gas composition and the Haardt &Madau UV ﬁeld).For Milky-Way sized galaxies accretion rates in the range 0.3-0.6h −1M ⊙yr −1(h =0.65)are found at z =0.One can show that the rate at which hot halo gas cools out is proportional to L X ,bol ·<1T>is the emissivity weighted inverse temperature of the hot gas.As mentioned previously we ﬁnd in the current work that L X ,bol is proportional to about the ﬁfth power of V c ,so asT ∝V 2c we would expect ˙M∝V 3c .This is indeed found by Sommer-Larsen et al.(2002)and is sensible since the I -band (and hence approximately mass)Tully-Fisher relation has a logarithmic slope of about three (Giovanelli et al.1997).Given the trends of L X with various environmental pa-rameters discussed in section 5.3we would expect galaxies formed in simulations with f b =0.05and primordial abun-dance to have about twice as large mass accretion rates,whereas galaxies formed in Z =1/3Z ⊙simulations will have accretion rates 3-4times lower than the similar primor-dial gas ones.Hence we expect at z =0a fairly strong trend of the ratio of present to average past accretion rate decreas-ing with increasing cooling eﬃciency.Indeed Sommer-Larsen et al.(2002)ﬁnd for their ΛCDM simulations accretion rates at z =1which are an order of magnitude larger than the ones at z =0,so disc galaxies may have been considerably more X-ray luminous in the past than they are today.A detailed analysis of mass accretion rates and high-z X-ray properties of our sample of simulated disc galaxies will be presented in a forthcoming paper.5.5Future X-ray observational testsFrom the predicted X-ray surface brightness proﬁles (Fig.4)and the halo temperatures (Fig.7)we have estimated the feasibility for detecting halo emission with XMM-Newton and Chandra using the most recent instrument responses.In order to avoid confusion with X-ray emission originating in the disc we aim at detecting halo emission from nearly edge-on disc galaxies at (vertical)disc heights of 10-15kpc.Count rates for a 5kpc high and 40kpc wide slice (parallel to the disc)at such disc heights were calculated assuming a column density of absorbing neutral hydrogen in the disc of the Milky Way of n H =2.5·1020cm −2(corresponding to the typical value for galactic latitudes of |b |∼60◦).For galaxies with circular speeds in excess of 300km s −1and distances d <∼50Mpc XMM-Newton,should be able to obtain a 5σde-tection at such disc heights in a 10ksec exposure.Although Chandra has a smaller collecting area than XMM-Newton the Chandra background is generally lower and its superior spatial resolution allows for more eﬃcient removal of con-taminating point sources.We thus expect that only slightly longer exposures are required for Chandra detection of halo emission than for XMM-Newton.However,the curve in Fig.4for the V c >300km s −1galaxies represents an optimistic case since the underly-ing simulations were run with primordial abundances and a strong external UV-ﬁeld,both increasing the present day X-ray luminosity.As mentioned in sections 5.3.2and 5.3.3,in more realistic simulations,including metals and a weaker external UV-ﬁeld,the halo ﬂux is lower by a factor of about 3.In this case,XMM-Newton as well as Chandra should still obtain a 5σdetection for V c >300km s −1galaxies within 25Mpc in about 25ksecs.For Milky Way sized galaxies,due to their much lower surface brightness (e.g.Fig.4)and lower halo temperature (the latter making these more sensitive to absorption)the predicted XMM-Newton and Chandra halo count rates are two orders of magnitude lower than for the V c >300km s −1galaxies.Detection of X-ray haloes for Milky Way sized galaxies at vertical disc heights of 10-15kpc will thus have to await future X-ray observatories with much larger collecting areas (Constellation-X and XEUS).6CONCLUSIONWe have presented X-ray properties of the hot gas haloes of disc galaxies derived from a large sample of physically realistic gravity/hydro simulations of galaxy formation and evolution.The simulated galaxies follow an L X,bol -V c relation with approximately the same slope as expected from simple cool-ing ﬂow models (L X,bol ∝V 5c ),but shifted to lower L X,bol ,c2001RAS,MNRAS 000,1–8。

Short GRBs and Mergers of Compact Objects Observational Constraints

SubjecΒιβλιοθήκη headings: gamma rays: bursts — ISM
1.
Introduction
Since data on GRBs started to accumulate over the past two decades, it was recognized that their time distribution appears to be bimodal6 , with about 25% of bursts having a short duration, of mean ∼ 0.2 sec, and the rest having a much longer duration, of mean ∼ 20 s (Mazets et al. 1981; Hurley et al. 1992; Kouveliotou et al. 1993; Norris et al. 2000). The separation between the two classes appears to be around 2 s. While the two classes of bursts seem to have similar (isotropic) spatial distributions, they diﬀer in several other respects. Short bursts tend to have harder spectra than long bursts (Kouveliotou et al. 1993; Dezelay et al. 1996), and about 20 times
Short GRBs and Mergers of Compact Objects: Observational Constraints

Segue

Sloan Extension for Galactic Underpinnings and Evolution(SEGUE)Segue(v.):To proceed to what follows without pauseHeidi Newberg1,Kurt Anderson2,3,Timothy Beers4,Jon Brinkmann3,Bing Chen5,Eva Grebel6, Jim Gunn7,Hugh Harris8,Greg Hennessy9,ˇZeljko Ivezic7,Jill Knapp7,Alexei Kniazev6,Steve Levine8,Robert Lupton7,David Martinez−Delgado6,Peregrine McGehee2,10,Dave Monet8,JeﬀMunn8,Michael Odenkirchen6,JeﬀPier8,Connie Rockosi11,Regina Schulte−Ladbeck12,J.Allyn Smith10,Paula Szkody11,Alan Uomoto13,Rosie Wyse13,Brian Yanny141Rensselaer Polytechnic Institute2New Mexico State University3Apache Point Observatory4Michigan State University5ESA/Vilspa,Madrid,Spain6Max-Planck-Institut f¨u r Astronomie,Heidelberg7Princeton University8US Naval Observatory,Flagstaﬀ9US Naval Observatory,DC10Los Alamos National Laboratory11University of Washington12University of Pittsburgh13Johns Hopkins University14Fermi National Accelerator LaboratoryI.Project SummaryA.Science GoalsWe propose an imaging and spectroscopic survey to unravel the structure,formation history, kinematics,dynamical evolution,chemical evolution,and dark matter distribution of the Milky Way.These results underpin our knowledge of the formation of the Milky Way Galaxy and of for-mation processes within the Milky Way,and will be the cornerstone of our understanding of galaxy formation processes in general.Cosmology is entering a precision era,facilitated by new work on the Cosmic Microwave Background by the Wilkinson Microwave Anisotropy Probe(WMAP)and on the distribution of galaxies by the Sloan Digital Sky Survey(SDSS)and by smaller,deeper large-telescope surveys.Galaxy formation and evolution,however,is still as data-starved as cosmology was twenty years ago.The Milky Way is the only galaxy in which we can hope to kinematically and spatially resolve the fossil record of evolution,since here the geometry is relatively unambiguous and intrinsically faint stars can be readily observed.Two key projects which focus on Galactic history and dynamics are:(1)detection of sub-structure in the Galactic halo,and(2)deﬁning and reﬁning our knowledge of the Milky Way’s Galactic disks.Other projects that will be necessary to realize the full potential of the key projects include understanding the relationship between SDSS photometry and the physical properties of stars,mapping the interstellar extinction in three dimensions,and studying the chemical evolution of the earliest Milky Way stars.To accomplish these goals,the existing SDSS hardware and software will be used to image (inﬁve SDSS passbands)four thousand square degrees of the Milky Way at low Galactic latitude and to sample the stars in this new region and in the existing high-latitude SDSS survey to targetand obtain spectra of240,000stars.These spectra will allow determinations of radial velocities and chemical abundances,which will allow us to study kinematics,dynamics,and the chemical history of the Galaxy.Key Project1:Halo SubstructureSDSS data have already been used to trace the tidal stream from the Sagittarius dwarf galaxy,and to discover a second large tidal stream in the plane of the Milky Way.The structure of the Milky Way’s halo is suﬃciently lumpy that it has so far deﬁed a consistent globalﬁt to the smooth component of the spheroid stars.The halo may contain a(possiblyﬂattened)power law component,aﬂattened inner halo,at least two large tidal streams,a dozen dwarf galaxies and a hundred globular clusters.Some of the dwarf galaxies and globular clusters are currently being pulled apart by tidal forces.Some of the stars in the components that appear to be smooth in density may retain kinematics of the stellar associations in which they were born.Disentangling the structure of the Milky Way halo requires that individual stellar associations can be separated by population(age and metallicity),kinematics,and spatial density.The SDSS and SEGUE photometry can be used for photometric parallax(spatial density)and isochrone ﬁtting to stellar components(into rough age and metallicity bins).Galaxy components can be separated kinematically with radial velocities.The stellar physical properties determined from spectroscopy and from imaging of open clusters serve as a check on the isochrones and further serve to illuminate the chemical composition of each component.The dynamics of stellar streams allow us toﬁt the Galactic potential’s shape and orientation, and constrain parameters characterizing the lumpiness of the halo dark matter.The dark matter itself could accrete with time as the stars do;knowledge of the Galactic merger history and grav-itational potential place important constraints on N-body models of galaxy formation and on the expected velocity distribution of dark matter particles.The velocity distribution of dark matter particles can aﬀect both the energy spectrum and annual modulation of the expected signal in dark matter direct detection experiments on the Earth.Key Project2:Disk structureImaging at low Galactic latitude will allow us to study the transition from thin disk to thick disk toﬂattened inner halo.There is general agreement on the number and exponential form of the Milky Way stellar disks,but little agreement on the exact parameters of each.This situation is similar to the state of our knowledge of the halo several years ago,when there was general agreement that the spheroidal population of stars was well modeled by a power law,but with no agreement on theﬂattening parameter(measurements range from c/a=0.4−1.0).As the current generation of precise data is beginning to show,the Milky Way’s disks are not as simple as the present models suggest.The fact that the stars associated with the large tidal stream in the plane of the Milky Way were initially widely attributed to an extended thick disk of stars underlines how little we know for sure about the number and detailed form of the Milky Way’s disks.SEGUE will use spectra(physical properties and radial velocities),photometry(stellar popu-lation and spatial density),and proper motions from outside catalogs to separate and normalize the Galactic components in the solar neighborhood.Additionally,the disk components will be traced using stars as close asﬁve or ten degrees from the Galactic plane,using techniques similar to those forﬁnding halo substructure.We do not expect to be able to trace all the spiral substructure of the young thin disk,as that is best studied at longer wavelengths,but we will be able to trace the structure,kinematics,and compositions of the other disks as a function of position in the Galaxy. This latter goal will require that we understand the three dimensional structure of the Interstellar Medium(ISM),including independent measures of dust from SEGUE observations,and its eﬀects on our stellar photometry.Formation HistoryThe identiﬁcation and characterization of the Milky Way components can be utilized as an archaeological“dig”illuminating the fossil record of galaxy evolution.We will study how many mergers and of what size and time series must have occurred to make the Milky Way.We will begin to be able to pick apart the chemical and dynamical evolution of the Galaxy as a whole.We will search for rare,especially low metallicity,stars that‘freeze-in’the state of the ISM during the earliest stages of star formation in the Universe.The rare stellar objects identiﬁed in this survey will provide followup targets of high scientiﬁc interest for the world’s largest telescopes.B.Survey DataApproximately four thousand square degrees of new imaging data,at moderate to low Galac-tic latitude,and spectra of240,000Galactic stars will be acquired.The imaging footprint was designed so that no part of the sky(aboveδ=−20◦observable from the Apache Point Observa-tory)is more than10◦from either an SDSS or a SEGUE imaging stripe.In the Galactic caps, no part of the sky is more than5◦from an imaging stripe.In addition,the scans are designed to tie the photometric calibration from the SDSS north Galactic cap to the scans in the south,and to cross each other a suﬃcient number of times to reduce systematic uncertainties in the overall photometric calibration.The positions of the spectroscopic plates are chosen to sample the Galaxy in all directions,so that no part of the observable sky is more than about ten degrees from a spec-troscopic plate,and to target well-studied open clusters.Figure1shows the approximate layout of the SEGUE imaging stripes and spectroscopic plates on the sky.The low Galactic latitude imaging enables studies of the metal-rich Galactic thin disk,the vertical structure of the thin and the thick disks,the Galactic warp andﬂaring,the three dimen-sional structure of the ISM,and present star forming regions.The imaging will be taken in similar weather conditions,at the same scanning rate(which translates to the same exposure time),and with the same instruments andﬁlters as the SDSS.Each stripe is2.5◦wide and requires two interleaved scans with the SDSS imaging camera,on separate nights,to complete.The imag-ing includes twelve constant Galactic longitude stripes which go through the Galactic plane and typically extend thirty-ﬁve degrees on either side(dashed green lines in Figure1).These stripes are separated by about20◦of Galactic longitude,varying somewhat to pass through known open clusters,SIRTF legacyﬁelds,and patches of low extinction near the Galactic plane.The constant Galactic longitude stripes connect and overlap SDSS imaging of the Galactic caps to facilitate the photometric calibration of both old and new data.In addition,three new SDSS stripes(72,79,and90,shown as solid green lines in Figure1) will be imaged in the South Galactic Cap.Only three stripes were imaged in the South Galactic Cap during the SDSS survey,and the additional stripes are needed to sample that part of the sky about every ten degrees.One stripe(red line in Figure1)follows the Sagittarius dwarf tidal stream across the northern sky.Two long(half-ﬁlled)strips(magenta lines in Figure1)cross the remaining SEGUE stripes,and will be used to cross-calibrate the stellar photometry to a level of at least2%(systematic+rms),and will trace low latitude structures,including the newly discovered tidal stream in the Galactic plane.The SDSS camera must scan along great circles,so all of these stripes describe great circles on the Celestial Sphere.The spectroscopic observations include1200spectra in each of200individual sky directions. The plate positions were chosen to sample the sky in all observable directions,and spectra will be selected to sample stars from one to100kpc from the Sun and from as many Galactic substructures as possible.Additional observations target regions of high interest such as open clusters,star formation regions,and known tidal streams in the halo.Each plate position is3◦in diameter.We will design two640-ﬁber plates in each plate position,with a total of about1200stellar targets-60-300306090120150180210240270300-90-60-300306090-90-60-30030609010230130l b s90s86s79s72s9s16s27s37s4418h22h2h 6h 6h 17.4h 2718.0h 618.6h 837.9h 498.6h 1621.1h -022.8h 25 1.3h 32 3.6h 17 5.0h -11Figure 1.Low Latitude Imaging and Spectroscopy Plan.The SFD (1998)reddening map is shown in Galactic coordinates;note the center is shifted to (l,b )=(120◦,0◦).Green,red,magenta (purple)lines indicate new SEGUE scans to be obtained.SEGUE Imaging at l =110degrees from −30◦<b <30◦has already been obtained as of Nov 2003(SDSS runs 4134,4135,4144,4152).The red line along the great circle with (node,incl)=(32◦,35◦)follows the Sagittarius dwarf tidal stream.Magenta lines indicate half-ﬁlled “strips”in portions of the sky at low Galactic latitude,and cross the constant longitude stripes for better calibration;the great circles arcs are deﬁned by (node,incl)=(259.9◦,43.6◦),(311.0◦,66.7◦).The total SEGUE imaging area is about 4000sq.degrees,of which 200sq.degrees has already been obtained.Black lines indicate existing or planned SDSS regular imaging.Black dotted reference lines are at b =0◦and declination (DEC)=−20◦(no SEGUE imaging or spectroscopy occurs at a DEC of <−20◦from Apache Point Observatory in the Northern hemisphere).Black long dashed lines mark constant Right Ascensions (RAs)of (18,22,2,and 6)bels above the top of the ﬁgure indicate RA,DEC start and end for a vertical SEGUE imaging stripe.Open black circles indicate positions of known Sagittarius dwarf tidal stream stars.Filled black circles indicate positions of known Monoceros stream stars.Open black diamonds indicate positions of known Galactic open clusters.The blue circles indicate individual 3-degree diameter positions of Galactic structure plate pairs (bright plate:45min exposure,plus faint plate:90min exposure),168blue plate pairs.Yellow circles indicate positions of ‘special plates,’landing on a known open cluster,the Sag.dwarf stream,or the Monoceros Ring structure.29blue plate pairs.Total:197plate pairs and about 240,000stellar spectra with resolving power R ∼2000,and 3<S/N <100for objects with 20.3>g >14.5.plus calibration objects.One plate will have the SDSS standard 45minute exposure time,and the other will be exposed for twice as long,allowing us to reach stellar targets as faint as g ∼20.3.Spectroscopic targets will include halo giants,metal-poor dwarfs,G disk and halo dwarfs,white dwarfs,and a large variety of rare stars.At low latitudes,targets within star-forming regions will be selected.II.Scientiﬁc CaseOur understanding of Galaxy evolution has advanced considerably since the monolithic col-lapse model of Eggen,Lynden-Bell and Sandage(1962)was adopted as the standard.Most experts now believe that the Galaxy was built up through a series of mergers(Searle&Zinn1978),though there is no agreement on the number and size of the merger events.These current models of galaxy formation stem from cold dark matter(CDM)simulations that show the outer halos accret-ing over billions of years(Steinmetz&Navarro2002),and from the increasing number of examples of moving groups and tidal disruption discovered in the halo of our galaxy(Majewski et al.2003; Newberg et al.2002;Odenkirchen et al.2001a;Irwin&Helmi et al.1999;Ibata,Gilmore,& Irwin1995;Grillmair et al.1995;Irwin&Hatzidimitriou1995),M31(Ferguson et al.2002),and other external galaxies(Shang et al.1998,Zucker et al.2004).It is possible that the hierarchical merging process is most important in the dark matter dominated galactic halos,while disks might form from the(angular momentum conserving)collapse of the gas within the stellar spheroid.However,Λcold dark matter simulations suggest merging may signiﬁcantly aﬀect the formation of disks as well(Abadi et al.2003).Evidence of mergers is currently most apparent in the outer halo where signatures of satellite accretion persist for many gigayears(Johnston,Spergel,&Hernquist1995).Dwarf galaxies and globular clusters are among the outer halo objects which are merging at the current epoch.It is also possible that there exist some lumps of dark matter in the outer halo that have not yet merged,and which do not contain stars(e.g.,Bullock,Kravtsov,&Weinberg2001).These latter structures could be evident by their perturbation of tidal tails and warping of disk structures.Our own Milky Way is the only galaxy that we can presently study at suﬃciently high spatial and kinematical resolution,and at suﬃcient depth,to address many of the open questions of galaxy formation and small-scale structure evolution in sub-halos.Our goal is to obtain the spectroscopic and photometric data required to unravel the structure,the formation history,the kinematic and dynamical evolution,the chemical evolution,and the distribution of the dark matter within and around the Milky Way.We propose two key projects,which contribute to our knowledge of the Galactic mass assem-bly and disk formation models.These projects are:(1)detection of substructure in the Galactic halo,and(2)deﬁning the structures of the Galactic disks.These are really two parts of the one key project to deﬁne the major components of the Milky Way galaxy,but are listed separately since they may require diﬀerent data sets and analysis methods.One may think of this proposed project as providing a large homogeneous input data set to a21st century model of the Galaxy–one which involves not only accurate multi-color photometry such as has gone into earlier models (Bahcall&Soneira1984),but large amounts of kinematic velocity and proper motion data which can be used to complete the dynamical and evolutionary picture.This technique is most similar to that used to construct the Besancon model of the Galaxy(Robin et al.2003),but with more input data.Detection of substructure in the Galactic halo requires photometry and radial velocities in as many directions as possible.The large tidal streams that have already been discovered are more than4kpc across,setting the scale over which we must sample the sky toﬁnd all large tidal structures.Characterizing the Galactic disks requires data collection primarily at low latitudes, and within a few kpc of the Galactic plane.The goal is to separate disk components by their stellar content,and then measure the global properties.These projects will separate and describe components using radial velocities,proper motions, chemical composition,photometric parallax,and isochroneﬁtting to photometry.In some sense, our survey is concentrated on the“big picture”of our galaxy.We will identify and constrain all of the largest components,paving the way for future inquiries which willﬁndﬁner substructure,S52-32-20.4S341+57-22.5S297+63-20.S6+41-20S223+20-19.4S200-24-19.8S167-54-21.5(RA,DEC)[l,b](0,0)[96,-60](15,0)[128,-63](30,0)[157,-58](45,0)[177,-49](60,0)[190,-37](75,0)[199,-25](90,0)[207,-11](105,0)[214,2](120,0)[221,15](135,0)[229,28](150,0)[239,41](165,0)[254,52](180,0)[276,60](195,0)[308,63](210,0)[337,58](225,0)[357,49](240,0)[10,37](255,0)[19,25](270,0)[27,11](285,0)[34,-2](300,0)[41,-15](315,0)[49,-28](330,0)[59,-41](345,0)[74,-52]15171921g *23171921231517192123151719212305101520253035404550Figure 2.Polar histogram of color-selected turnoﬀstars on the celestial equator .This data was obtained with the Sloan Digital Sky Survey.The radial distance gives the apparent dereddened g ∗magnitude.The angular position gives the RA.Galactic coordinates are labeled for reference.The shading of each cell indicates the relative number of counts of stars in each (r,θ)=(g ∗,α)bin.A typical absolute magnitude for stars with these colors is M g ∗=4.2.The feature at α=60◦is an artifact from imperfect reddening correction of a large dust cloud at this position.The streak at α=229◦represents stars in the globular cluster Palomar 5.The boldface labels indicate our names for identiﬁed structures of halo stars.S341+57-22.5at g ∗∼22.5,and S167-54-21.5at g ∗∼21.5are cross sections of the Sagittarius dwarf galaxy tidal stream.more accurately determine the chemical evolution,and measure proper motions for a large fraction of the stars in the Milky Way.A calibrated catalog of images,spectra,and associated derived quantities,will be the primary product of this survey.These will be generated in nearly real time,to be used for rapid follow-up work or as input targets to space-based or large aperture telescopes.The need for this global picture of our galaxy is well illustrated by the results of Newberg et al.(2000;Figure 2).In this ﬁgure,there are seven marked concentrations of stars.Concentrations S341+57-22.5and S167-54-21.5have been identiﬁed as cross sections through the tidal stream ofthe Sagittarius dwarf spheroidal galaxy,which is currently in the process of tidal disruption.The overdensities S223+20-19.4and S200-24-19.8are thought to be pieces of another tidal disrupting stream in the plane of the Milky Way galaxy.The concentrations in this region near the Galactic plane,at15th and17th magnitude near the anticenter are not named,but also are not understood in any global picture of the Milky Way.The overdensity S297+63-20.is thought to be another tidal stream,possibly associated with the Sagittarius dwarf galaxy,though this has not been conﬁrmed and remains controversial.The concentrations S6+41-20and S52-32-20.4are thought to be portions of the stellar spheroid,though their density distributions do notﬁt standard spheroidal models within the errors of our density measurements.One sees in Figure2a strong argument for a global view of the whole Milky Way,including low Galactic latitudes,since one cannot identify substructure without understanding the major Galactic components in which that substructure is embedded,and properly accounting for inter-stellar extinction.Many of the scientiﬁc analyses that we anticipate will be based on these data have counterparts in the much smaller-scale eﬀorts of individuals or groups,which,unavoidably, dilute their impact by acquiring data in a piecemeal and non-uniform fashion.A uniform survey creates a synergy which allows more global questions to be addressed and leaves behind a legacy data set against which future data sets will be compared.A better understanding of the Milky Way’s structure and evolution is already a“cornerstone”project in ESA’s science planning,through the GAIA satellite mission,and plays an important role in the deﬁnition of the science goals for NASA’s SIM and TPF missions,which are“key elements in NASA’s Origins Program.”We demonstrate here how a deep imaging and spectroscopic optical survey will complement as well as lay the groundwork for these ambitious satellite projects. Furthermore,SEGUE willﬁll a unique and vital niche complementing ongoing and planned large ground based Galactic structure programs such as RAVE and K.A.O.S.A.Characterization of the Component Stellar Structures in the Milky WayThe fossil record of galaxy evolution(star formation and mass assembly)is written in the chemical,kinematic,and spatial distribution of Galactic stars.The main recognized components of the Galaxy are the thin disk,the thick disk,the bulge,and the stellar spheroid.Recently many groups of astronomers have identiﬁed examples of Galactic structure that either requires additional components or an increase in the complexity of the traditional components.Kinematic studies show the existence of moving groups and coherent streams(numerous studies),and a group of stars(Gilmore,Wyse,&Norris2002)that may be part of the merger that puﬀed up the thick disk.Overdensities of stars over the Galactic bar(Parker,Humphreys,&Larsen2002)have been found in photometry.Also,a new metal-weak thick disk component has been proposed(see Norris 1994,Beers et al.2002and references therein).Clearly,even the basic stellar components of the Milky Way are not yet understood in depth. The complex substructure now being identiﬁed has undoubtedly biased the previous limited studies of thick disk structure,contributing to our present imprecise knowledge of the thick disk;study of many lines of sight over much of the sky will be necessary to unravel the substructure and obtain a more complete picture.This survey would speciﬁcally target the thick disk/halo boundary and substructure.The structures would be studied in stellar density from statistical photometric parallaxes,and in kine-matics through statistics of the radial velocities/metallicities in each component.We use the term ”statistical photometric parallax”to describe the method of using photometry to determine dis-tance(photometric parallax)in cases where the number of stars is large,so that statistics can be used to estimate the underlying spatial structure of the group.We will have a unique opportunity to study the stellar Metallicity Distribution Function(MDF),especially in the region where the thick disk and spheroid populations overlap.Figure3demonstrates preliminary results from SDSS-4-3-2-101[Fe/H]100010000100000Z D i s t a n c e (k p c )Figure 3.The distance and metallicity distribution of EDR stars.The distance distribu-tion of ∼4000stars from the SDSS Early Data Release (EDR)as a function of metallicity [Fe /H].One can clearly discriminate the presence of thick-disk stars with metallicities in the range -1.0<[Fe/H]<0and locations within a few kpc of the Sun,from the halo objects at large distances that extend to much lower metallicities.EDR spectra.Flaring and Warping of the Galactic PlaneThe disks of many galaxies both ﬂare and warp in their outer regions.Flaring is attributed to an increasing ratio of spherically distributed dark mass to disk mass with increasing distance from the center of the galaxy,and provides one of the few available methods of measuring the three-dimensional distribution of dark matter within a galaxy.The origin of Galactic warps is still something of a mystery.Tidal interactions with satellites and neighbors is an obvious cause;for example,the warp in the Galaxy is often attributed to the tidal inﬂuence of the Magellanic Clouds (e.g.,Weinberg and Nikolaev 2001;Garcia-Ruiz et al.2002)or of the Sagittarius dwarf spheroidal galaxy (Bailin 2003).However,not all warped galaxies appear to have (presently detected)neigh-bors.The warp andﬂare in the outer Galactic disk has been studied primarily using(radio)obser-vations of neutral hydrogen.The depth and color sensitivity of SDSS will allow the3-D structure of the northern warp in the Galaxy to be traced to distances beyond the entire known extent of the warp(20kpc,Binney1992),using photometrically identiﬁed giants,carbon stars(especially when combined with2MASS data)and red clump stars(e.g.,Margon et al.2002;Helmi et al.2003). The Structure of the Thick DiskHawley et al.(2002)show that early M dwarfs can be traced to distances of up to1kpc above the Galactic plane,well into the domain where the thick disk population dominates that of the thin disk.Since the stars of the thick disk are more metal-deﬁcient(typically by at least0.3-0.5 dex)than the thin disk stars,their colors,especially g−r,diverge from those of the metal-rich disk stars.One can therefore,at least in a statistical sense,separate the two populations.Aﬁrst look at the vertical structure of the thick disk from SDSS data has been carried out by Chen et al.(2001). Star counts in the thin and thick disks can be used to determine the initial mass function,and in particular the counts of lower-metallicity stars must be understood(see the recent discussion by Zheng et al.2001and Chabrier2003).The current SDSS imaging data provide star counts at high Galactic latitude only;under-standing and disentangling the vertical structure of both the thick and thin disks requires data covering the whole range of Galactic latitude at many longitudes.The Structure of the Galactic HaloSDSS data have already demonstrated the presence of very large structures in the Galactic halo(Yanny et al.2000;Ivezic et al.2000;Odenkirchen et al.2001a,b;Newberg et al.2002; Rave et al.2003;Yanny et al.2003).These structures include extra-tidal features around globular clusters and vast comoving stellar streams from accreted dwarf galaxies.These structures,and similar more tenuous analogues,may cover a signiﬁcant fraction of the sky.Imaging more sky allows such structures to be traced to larger angular sizes,and allows structures which do not completelyﬁll a great circle on the sky to be detected.Furthermore,high-metallicity globular clusters tend to be found at lower Galactic latitudes;there may be streamers and tails of diﬀerent color(metallicity)at lower latitudes.Indeed,recently Frinchaboy et al.2004showed that many low latitude open and globular star clusters are likely to be associated with a single large tidal stream in the Galactic plane.There is increasing evidence that even“globally recognized”structures,such as“the halo,”change dramatically with increasing Galactocentric distance.An inner,“ﬂattened”halo compo-nent,for example,has been indicated by many recent studies(Lemon et al.2003,Chiba&Beers 2000).Preston,Shectman&Beers(1991)have used the mean colors of blue horizontal-branch stars to indicate a possible decrease in the ages of stars with increasing Galactocentric distance. Sirko et al.(2004)use an analysis of the spectra of high-latitude A stars to demonstrate a change in the velocity ellipsoid of the halo with distance from the Galactic center.Both of these represent key results for understanding the formation of the Milky Way,which could be readily reﬁned and extended using our proposed survey eﬀort.Complexity of Galactic spheroid populations in other galaxies is traced by their globular cluster systems(at least in early-type galaxies with populous cluster systems that can be studied). Typically the globular cluster systems are bimodal,with bluer(lower metallicity)clusters mak-ing up a more extended and dynamically hotter system,and redder(higher metallicity)clusters comprising a system more centrally concentrated and sometimes rotating(e.g.many papers in IAU Symposium207,ed.Geisler,Grebel&Minniti2002).Typically theﬁeld-star populations have color and spatial distributions suggesting properties more similar to(and perhaps having。

A 1.4-GHz Arecibo Survey for Pulsars in Globular Clusters

a r X i v :0707.1602v 1 [a s t r o -p h ] 11 J u l 2007A 1.4-GHz Arecibo Survey for Pulsars in Globular ClustersJ.W.T.Hessels 1,∗,S.M.Ransom 2,I.H.Stairs 3,V.M.Kaspi 1,and P.C.C.Freire 41Department of Physics,McGill University,Montreal,QC H3A 2T8,Canada;hessels@physics.mcgill.ca2National Radio Astronomy Observatory,520Edgemont Road,Charlottesville,VA 229033Department of Physics and Astronomy,University of British Columbia,6224Agricultural Road,Vancouver,BC V6T 1Z1,Canada 4NAIC,Arecibo Observatory,HC03,Box 53995,Arecibo,PR 00612∗Current Address:Astronomical Institute “Anton Pannekoek”,University of Amsterdam,Kruislaan 403,1098SJ Amsterdam,The Netherlands ABSTRACT We have surveyed all 22known Galactic globular clusters observable with the Arecibo radio telescope and within 70kpc of the Sun for radio pulsations at ∼1.4GHz.Data were taken with the Wideband Arecibo Pulsar Processor,which provided the large bandwidth and high time and frequency resolution needed to detect fast-spinning,faint pulsars.We have also employed advanced search techniques to maintain sensitivity to short orbital period binaries.These searches have discovered 11new millisecond pulsars and 2promising candidates in 5clusters,almost doubling the population of pulsars in the Arecibo-visible globular clusters.Ten of these new pulsars are in binary systems,and 3are eclipsing.This survey has discovered signiﬁcantly more very fast-spinning pulsars(P spin 4ms)and short orbital period systems (P orb 6hr)than previoussurveys of the same clusters.We discuss some properties of these systems,aswell as some characteristics of the globular cluster pulsar population in general,particularly its luminosity distribution.Subject headings:globular clusters:general —pulsars:general —binaries:gen-eral —radio continuum:stars —stars:neutron1.IntroductionCurrently,there are approximately130known pulsars in globular clusters(GCs)1,of which about60%are observed to be in binary systems.2Roughly two thirds of all the pulsars known in GCs have been discovered in only the last seven years by surveys(e.g. Possenti et al.2003;Ransom et al.2005a;Hessels et al.2006)using low-temperature re-ceivers(T rec<∼35K)at central observing frequencies betweenνcenter=1−2GHz,large bandwidth and high time and frequency-resolution backends(e.g.the Wideband Arecibo Pulsar Processor,Dowd,Sisk,&Hagen2000,and the GBT Pulsar Spigot,Kaplan et al. 2005),advanced search techniques for binaries(e.g.Ransom,Eikenberry,&Middleditch 2002;Ransom,Cordes,&Eikenberry2003;Chandler2003),and copious amounts of pro-cessing time on dedicated computer clusters.Though a few non-recycled pulsars have been found in GCs(e.g.PSR B1718−19in NGC6342,Lyne et al.1993),almost all known GC pulsars are MSPs.In fact,the150known GCs3orbiting the Milky Way contain roughly three orders of magnitude more observed millisecond pulsars(MSPs)per unit mass than the Galactic plane,which contains approximately60known MSPs.GCs have proven to be the most fruitful place to look for MSPs partly because of the enormous stellar densities in their cores,which exceed those in the Galactic plane by up to six orders of magnitude.These conditions promote many diﬀerent formation processes that create binary systems in which a neutron star can be spun-up,or“recycled”(Alpar et al.1982;Radhakrishnan&Srinivasan 1982),through the accretion of matter from its companion star(see Camilo&Rasio2005, the most up-to-date general overview of pulsars in GCs,for a review of the formation and evolutionary processes at work in the cores of GCs).Furthermore,the cores of GCs,where most of the MSPs reside,typically have radii less than an arcminute,small enough to be covered by a single telescope pointing.This aﬀords the possibility of making single,deep multi-hour integrations for multiple faint MSPs in GCs,something that is not feasible in large area surveys of theﬁeld.Some clusters contain many pulsars:Terzan5and47Tucanae harbor33and22known pulsars,respectively.Together,they contain43%of the total known GC pulsar population (Camilo et al.2000;Ransom et al.2005a;Hessels et al.2006).Finding numerous pulsars in a single cluster allows interesting studies of the cluster itself,in addition to the individualpulsars contained therein.Such studies have included the detection of intra-cluster ionized gas(Freire et al.2001),high mass-to-light ratios(D’Amico et al.2002),and hints at the cluster’s dynamical history(Possenti et al.2003).However,many clusters still contain no known pulsars at all,despite sensitive searches(Camilo&Rasio2005).In some cases,this may simply be because pulsars are generally intrinsically weak objects and GCs are often distant(approximately90%of the GCs orbiting the Milky Way are>5kpc from the Sun). Interstellar scattering,which broadens pulsations because of multi-path propagation,can also be a major obstacle,especially for MSPs in clusters at low Galactic latitudes.At low radio frequencies( 500MHz),scattering can completely wash out the signal of fast-spinning pulsars.While GCs are clearly the most proﬁtable targets forﬁnding MSPs,these searches remain non-trivial.First,the requisite short sampling times(<100µs),high frequency resolution(<0.5MHz channels),long integrations(few hours),and large bandwidth(> 100MHz)of these data can make data acquisition and storage formidable(data rates∼30−100GB/hr).Recent surveys with relatively good sensitivity to binary pulsars have re-vealed that the majority of GC MSPs are in binaries.The pulsed signal from these binary pulsars is smeared in Fourier space by their orbital motion,and thus even a bright pul-sar can go undetected if nothing is done to correct for this modulation over the course of a long observation.Advanced and computationally intensive techniques,which partially search orbital parameter space,are required to recover the majority of the lost signal.The extra eﬀort required toﬁnd binaries is well justiﬁed however,as some of the most exotic pulsar binaries known have been found in GCs.For instance,PSR B1620−26in the cluster M4is in a hierarchical triple system with a white dwarf and a1−3M Jup planet,the only planet known in a GC(Thorsett et al.1999;Sigurdsson&Thorsett2005);PSR J0514−4002 in NGC1851,with an eccentricity of0.89,is one of the most eccentric binary pulsars known(Freire et al.2004);PSR B2127+11C in M15is a rare double neutron star binary (Anderson et al.1990);PSR J1748−2446ad in Terzan5is the fastest-spinning neutron star known(Hessels et al.2006);and a few MSPs with possible main-sequence companions have been found(e.g.PSR J1740−5340,D’Amico et al.2001,similar systems have not been found in the Galactic plane).Other exotic binaries,perhaps even an MSP-MSP binary or a MSP-black-hole binary,may have eﬀective formation channels only in GCs(but see Sipior,Portegies Zwart,&Nelemans2004).Here we present searches for radio pulsations from22GCs,using Arecibo at central ob-serving frequencies between1.1−1.6GHz.Roughly half of these clusters have been searched previously with Arecibo at430MHz(Anderson1993;Wolszczan et al.1989a,b).The highly successful Anderson(1993)survey found11of the15pulsars known in these clusters prior to the survey presented here.Our searches have uncovered11new MSPs and2promisingcandidates inﬁve clusters:M3,M5,M13,M71,and NGC6749.Acceleration searches were crucial inﬁnding all but2of these systems.Ten of the new pulsars are in binaries,3of which are eclipsing,with orbital periods of only a few hours.For comparison,no eclipsing systems were known in these clusters prior to our survey and only5of the15previously known pulsars were in binaries.Of the pulsars presented here,9out of11have P spin<4ms. Only2of the15previously known pulsars in these clusters have P spin<4ms,clearly demon-strating the improved sensitivity of this survey over past surveys to the fastest MSPs.In§2 we describe the targets,observational setup,and sensitivity of the survey.In§3we outline the search procedure and analysis pipeline.In§4we present the results of the survey.In the discussion of§5,we comment on the characteristics of the GC pulsar population in general, particularly its luminosity distribution.In§6,we conclude.2.Observations2.1.TargetsWe observed every known Galactic GC visible from Arecibo4and within70kpc of the Sun without any selection bias towards larger or denser clusters.The sample of22GCs is listed in Table1,along with basic and derived cluster parameters(Harris1996,unless otherwise indicated,all GC quantities used in this paper are from the2003February revision of the catalog).The numbers of known isolated and binary pulsars in each cluster are indicated,withﬁgures in parentheses denoting the number of pulsars found by this survey. For clusters containing known pulsars,the average dispersion measure(DM)of the pulsars is indicated,as well as their spread in DM,which is in parentheses(for clusters where two or more pulsars have been found).For clusters with no known pulsar,the DM is also unknown;the values listed are derived from the Cordes&Lazio(2002)“NE2001”model for the distribution of free electrons in the Galaxy,using the Harris(1996)position and distance to the cluster.2.2.Data AcquisitionEach cluster was observed at least twice for the full time it is visible with Arecibo (Table1).The clusters were observed in one of two campaigns in the summers of2001and2002using the Gregorian L-band Wide receiver5(T rec∼35K).Depending on the known or predicted DM of the cluster,the central observing frequency was either1175MHz(DM 100pc cm−3)or1475MHz(DM 100pc cm−3).Our observations were made using the Wideband Arecibo Pulsar Processor(WAPP,see Dowd et al.2000,for details),a digital auto-correlator with conﬁgurable sampling time(3or9-level samples)and number of lags. Generally,3-level samples were autocorrelated with256lags,accumulated every64-µs,and summed in polarization before being written to the WAPP disk array as16-bit numbers. For the few clusters with known or predicted DMs greater than100pc cm−3,we used128-µs sampling and512lags.These conﬁgurations were chosen in order to minimize dispersive pulse smearing and to take full advantage of the WAPP’s maximum sustainable data rate at the time of the observations(8MB/s,or∼30GB/hr).Data were transferred to DLT magnetic tape for oﬄine analysis and archiving.These observations resulted in about4TB of data on about100tapes.At the time of our original cluster search observations,only one WAPP backend was available,providing100MHz of bandwidth.In more recent observations of M3,M5,M13, M71,and NGC6749(after December2002),which were made as part of the timing observa-tions of the new discoveries,we used three of the four available WAPP backends,each with 100MHz of bandwidth centered at1170,1420and1520MHz.The frequency gap between the lower and upper bands was to avoid persistent and intense radio frequency interference (RFI)in the frequency ranges1220−1360MHz and>ing the PRESTO6pul-sar software suite,these data were partially dedispersed into a reduced number of subbands (generally16or32subbands)at the average DM of the cluster pulsars before further tim-ing or search analysis.This process aﬀords an order-of-magnitude reduction in data size (enabling transfer of these data over the internet from Puerto Rico to Canada),while still providing the possibility of creating dedispersed time series at a variety of DMs around the average DM of the cluster.These data were also searched for new pulsars,in the manner described in§3.2.3.Search SensitivityWe can estimate the typical minimumﬂux density to which our searches were sensitive, as a function of the radiometer noise and observed pulsar duty cycle,using the equationS min=σξT sysn∆νT obsw obssearching or folding,t∆DM,and interstellar scattering,t scatt.One can express the observed width as the sum in quadrature of these terms:w obs2=w int2+t samp2+t DM2+t∆DM2+t scatt2,(2) where the dispersive smearing(assuming the channel bandwidth∆νchan≪νcenter)across an individual channel is given byt DM=8.3 DM MHz νcenterGHz −2− νhigh pc cm−3 ms,(4)(whereνlow andνhigh are the low and high frequency edges of the bandwidth respectively), and the scattering(t scatt is in ms)can be estimated by the empirical formula(Bhat et al. 2004)log10(t scatt)=−6.46+0.154log10(DM)+1.07(log10(DM))2−3.86log10 νcenterDM and period.7We compare the sensitivity of our survey with that of the only other major survey of these clusters with Arecibo(Anderson1993).As the Anderson(1993)survey was conducted at430MHz,we use a typical pulsar spectral index of−1.8(Maron et al.2000) as well as aﬂatter spectral index of−1.3to scale that survey to1400MHz,so that the sensitivities of both surveys can be more directly compared.8For low-DM pulsars with spin periods 10ms and a standard spectral index(α=−1.8),the Anderson(1993)survey had a similar sensitivity to ours.However,for pulsars spinning faster than this,especially those at high DMs and with relativelyﬂat spectral indices,our survey provides a signiﬁcant increase in sensitivity.For example,we were roughly3−6times more sensitive to a2-ms pulsar at a DM of100pc cm−3,assuming−1.8<α<−1.3.Our searches are likely the deepest searches for GC pulsars yet undertaken.Though recent searches with Parkes(e.g.Camilo et al.2000)and GBT(e.g.Ransom et al.2005a) beneﬁt from larger recordable bandwidth and longer possible source tracking times,Arecibo’s much larger gain still more than compensates for these(Table2).Comparing raw gains, maximum tracking times,and recordable bandwidths,simple scaling using Equation1and a spectral index of−1.8to compare diﬀerent observing frequencies,shows that,with the currently available data-recorders,Arecibo has a raw sensitivity4×greater than Parkes and 2×greater than the GBT for an isolated pulsar.For binary pulsars,where blind search sensitivity doesn’t improve as T1/2(see§3.2and5.1.1.),and longer tracking times provideobsless beneﬁt,Arecibo is more sensitive by an additional factor of roughly2.3.Analysis3.1.Radio Frequency Interference ExcisionPeriodicities due to terrestrial radio frequency interference(RFI)can swamp search candidate lists.Ultimately this reduces a survey’s sensitivity by increasing the number of false positives.First,strong bursts of interference,principally from airport and military radars,were removed in the time domain by clipping samples found to be further than6σfrom the mean in the DM=0pc cm−3time series.Secondly,we created time-frequencymasks and applied them to all data before searching.These masks remove certain channelsduring speciﬁc time intervals when either the maximum Fourier power,standard deviation,or mean of the data surpass statistically determined thresholds.This method was veryuseful at excising strong narrow-band and transient RFI from the data before searching and typically only<∼10%of the data were stly,we removed known“birdies”(weak buthighly periodic broadband interference and bright known pulsars)from the power spectrumby setting the powers in these narrow frequency intervals to zero.We found that using these techniques and observing in the relatively RFI-clean frequency ranges of1120−1220MHzand1370−1570MHz kept the size of our candidate lists manageable,while minimizing therisk of rejecting potentially interesting candidates.3.2.Search TechniquesIn a standard Fourier-based search,which identiﬁes pulsar harmonics above a certain threshold,binary pulsars are strongly selected against because of orbital modulation,which smears the signal power over many bins in frequency space.As the majority of MSPs inGCs are in binary systems,sometimes with orbital periods comparable to or shorter thanthe observation length,it is crucial to use more sophisticated techniques.Our primary search method is a Fourier-based matchedﬁltering technique(Ransom et al. 2002)which assumes that the pulsar’s orbit can be described by a constant acceleration(i.e.constant frequency derivative)over the course of the observation.This technique is a frequency-domain version of a previously used constant acceleration technique(Middleditch&Kristian 1984;Anderson et al.1990;Camilo et al.2000)in which the time series is quadratically stretched or compressed before it is Fourier transformed,in order to simulate diﬀerent accel-erations.The advantage of the Fourier-based technique used here is that it is computationallymore eﬃcient than the equivalent time-domain-based method,and provides even samplingof frequency derivative space.This method is typically most sensitive to binaries where theorbital period is>∼10×the integration time of the observation,although very bright pulsarswith orbital periods much shorter than this can often be detected(Johnston&Kulkarni 1991).As our Arecibo observations were typically from1−2.5hr in length,for full tracksthis search method is most sensitive to orbital periods>∼10−ing the same tech-nique,we also searched overlapping subsections of the data sets(which were typically aboutone third the total observation length),in order to be sensitive to larger accelerations andtighter binaries(P orb 3−10hr),as well as eclipsing systems.To search for binaries with orbital periods a factor of∼2or more shorter than theobservation time,we used a“phase modulation”search(Ransom et al.2003;Jouteux et al. 2002).This method makes use of the change in phase of the pulsations from a binary pulsar and the“comb”pattern created by such a signal in the Fourier power spectrum.By taking Fast Fourier Transforms(FFTs)of short subsections of the power spectrum from the full observation,one can detect the sidebands,which are evenly spaced at the orbital period and centered on the true frequency of the pulsar.This method is appropriate for binary pulsars with orbital periods less than half the observation length,and increases in sensitivity as the ratio T obs/P orb increases.Given the1−2.5hr integration times used in this survey, this technique was sensitive to an area of orbital parameter space in which no binary radio MSPs have yet been observed.3.3.Data Analysis PipelineHere we summarize the data analysis pipeline applied to the observations of each cluster. The routines used in the pipeline are part of the PRESTO pulsar search and analysis pack-age.First,we generated an RFI mask from the raw data(as described in§3.1).Applying this mask,a DM=0pc cm−3,topocentric time series was created and the corresponding power spectrum was examined by eye for periodic interference.We added the worst of the interference to the“birdie list”of frequency intervals to be subsequently removed from the Fourier transform.Again applying the RFI mask,we then dedispersed the data at a num-ber of trial DMs using DM steps≤1pc cm−3,clipping samples found to be greater than 6σfrom the mean in the DM=0pc cm−3time series.The time series were transformed to the solar-system barycenter during dedispersion,which allowed us to compare directly candidate periods from separate observations.This was very useful for distinguishing likely pulsar periodicities from the RFI background,which can vary greatly between observations. For clusters with known pulsars,we dedispersed at a range of trial DMs equal to at least 10%of the cluster’s average DM.For clusters with no known pulsars,we created time series at a range of DMs centered around the predicted DM of the cluster(Cordes&Lazio2002), given its Galactic coordinates and distance in the Harris(1996)catalog.We assumed at least a100%error in the predicted DM value when choosing a DM search range.For example, for a cluster with a predicted DM of100pc cm−3,we searched dedispersed time series with DM=0−200pc cm−3.The dedispersion and subsequent analysis of the time series was con-ducted in parallel using multiple processors on a52-node dual-processor Linux cluster called “The Borg”,located at McGill University and constructed by our group speciﬁcally for pul-sar searches.Once the dedispersion was complete,the time series were Fourier-transformed. Our analysis did not restrict the number of samples in the time series to be a power oftwo.9We then set the frequency intervals of the birdie list to zero in the power spectra before searching the power spectra with both the phase-modulation and matched-ﬁltering techniques described in§3.2.As most pulsars have relatively short duty cycles,with spectral power divided between numerous harmonics,we summed harmonics in our matched-ﬁlter search to increase sen-sitivity to such signals.For each candidate signal,sums of1,2,4,and8harmonics were tried and the optimum combination was determined.In these searches,we looked for signals where the highest of the harmonics used in summing drifted by up to z max=170bins in the Fourier domain.Higher order harmonics will drift by N harm times more bins than the fun-damental,where N harm=1for the fundamental,N harm=2for the second harmonic,etc.If, for example,8harmonics were summed in the identiﬁcation of a particular candidate period, then the maximum number of bins the fundamental could drift by during the observation and still be detectable by our search would be170/8.Conversely,for signals where the fundamental drifted by more than170/2bins during the observation,we were sensitive to at most one harmonic.The maximum number of bins a signal is allowed to drift corresponds toa maximum line-of-sight acceleration of a max=z max cP spin T−2obs N−1harm≃4P spin,ms T−2obs,h N−1harm ms−2,where P spin,ms and T obs,h are the spin period in milliseconds and the observation time in hours and N harm is the highest order harmonic used in summing.Overlapping subsections of the observation,corresponding to about a third of the total observation length,were also searched using the matched-ﬁltering technique and z max=170bins,in order to look for more highly accelerated pulsars.The resulting candidate lists from the matched-ﬁltering search were generally short enough that they could easily be examined by eye,although we also parsed candidate lists with a script that automatically folded candidates above an equivalent gaussian signiﬁcance threshold of7.Interesting candidates were folded using the estimated DM,period,and period derivative from the search and these parameters were optimized by the folding software to maximize the S/N ratio of the folded proﬁle.The output plot from such a fold was used to determine whether a given candidate warranted further attention(see sample discovery plot and description in Figure2).To identify potentially interesting candidates from the phase-modulation searches,we compared the search outputs from the diﬀerent observing epochs of each cluster.The criteria for a promising phase-modulation candidate was a signal with a signiﬁcance>10σthat did not peak in signiﬁcance at DM=0pc cm−3,had an orbital period>600s,and appeared,by virtue of a similar orbital period,in the candidate lists of at least two observing epochs.4.Results4.1.RedetectionsTable3lists the15previously known pulsars in our survey clusters.The majority of these pulsars were easily detected by our search pipeline(Table3).In fact,because many of these sources are relatively bright,masking these periodicities and their many signiﬁcant harmonics was an important factor in reducing the length of candidate lists.The only previously known pulsars not detected in our searches are M15F,G,and H.10Although these are all isolated pulsars,and are well within the∼1.5′half-power radius of the1.4-GHz Arecibo main beam,our non-detections are not very surprising:these are the dimmest pulsars in M15,and were all found in searches of multiple,combined observing epochs (Anderson1993).Assuming standard spectral indices(α=−1.8),they haveﬂux densities right at the limit of our search sensitivity.4.2.New MSPsWe have discovered11MSPs and2promising candidates in5clusters.Three of the clusters with new pulsars(M3,M71,and NGC6749)contained no known pulsars prior to our survey.All of the pulsars discovered in this survey were found using the matched-ﬁltering acceleration search technique,and all but the newly-found isolated pulsar M13C (PSR J1641+3627C)and the long-orbital-period binary M3D(PSR J1342+2822D)required a non-zero trial frequency derivative(acceleration)in order to be detectable.Given the criteria outlined in§3.3,no interesting candidates were identiﬁed by the phase-modulation search.The spin periods of the new pulsars have a narrow range2.4−5.4ms.For comparison, the previously known pulsars in these clusters have spin periods in the range3.5−111ms, with only2pulsars having P spin<4ms.All but one of the new pulsars is in a binary system, with orbital periods ranging from2.1hr up to129d.Although none of the previously known binaries in these GCs show eclipses,3of the new pulsars found here eclipse.The basic characteristics of these pulsars are summarized in Table4.Integrated pulse proﬁles,which are often the sum of numerous observations,are shown in Figure3.We have conducted monthly timing observations of these discoveries over the course of approximately two years using Arecibo and multiple WAPP backends.In fact,M3D (PSR J1342+2822D),M5E(PSR J1518+0204E),and M13E(PSR J1641+3627E)were alldiscovered in searches of timing data because of fortuitous scintillation.The timing resultsfor the M5and M71pulsars will be presented by Stairs et al.(2008,in preparation)and thetiming of the M13pulsars will be presented by Ransom et al.(2008,in preparation).4.2.1.M3(NGC5272)We have found theﬁrst three,and likely four,pulsars known in M3.All of thesepulsars were detected in observations in which interstellar scintillation increased theirﬂuxto a detectable level and they are not consistently detectable with Arecibo.They have DMswithin0.3pc cm−3of each other,with an average DM of26.4pc cm−3.This comparesvery well with the23pc cm−3DM predicted by the NE2001model(Cordes&Lazio2002).On these grounds alone,there is little doubt that these pulsars are members of M3.M3A(PSR J1342+2822A)is a2.54-ms binary which has only been detected three times,on MJDs52491,52492and52770(we list the number of detections of each new pulsar in Table4).Due to this paucity of detections,we do not currently know the orbital parameters for thispulsar,although its orbital period is likely on the order of a day.We also cannot derivea precise position for this pulsar,though we note that,since the half-power radius of theArecibo1.4-GHz main beam is∼1.5′,it is likely no further than this from the cluster center.M3B(PSR J1342+2822B)is a2.39-ms binary in a34.0-hr orbit with a0.2M⊙(minimummass)companion,which,in analogy to other systems with similar orbital parameters,maybe a low-mass helium white dwarf.It is the most consistently detectable pulsar in M3,and isvisible at least faintly in roughly half of our observations.On some occasions M3B has shownvery large increases inﬂux(up to aﬂux density at1400MHz S1400∼0.1mJy),presumably due to diﬀractive scintillation.We have derived a phase-connected timing solution for M3B(Table5)using the TEMPO pulsar timing package11and standard pulsar timing techniques.This solution places it8.3′′(0.25core radii)from the center of the ing thisprojected position and the simple cluster model outlined in Freire et al.(2005),we canderive the maximum contribution to the observed period derivative from acceleration in thecluster’s gravitational potential.In this model,weﬁnd|a max/c|P spin=1.6×10−20s/s,whichis comparable to M3B’s observed period derivative.From this we place an upper limit onM3B’s intrinsic period derivative and corresponding lower limits on its characteristic ageand dipole magneticﬁeld(Table5).M3D has a spin period of5.44ms and an orbital period∼129d.Though the detectionsof M3D are too sparse to derive a phase-connected timing solution for this pulsar,by inserting。

长波敏化发光铕配合物纳米粒子的制备与表征

[Article]物理化学学报(Wuli Huaxue Xuebao )Acta Phys.-Chim.Sin .,2010,26(7)：2031-2036July Received:February 2,2010;Revised:March 27,2010;Published on Web:May 21,2010.*Corresponding author.Email:wangy@;Tel:+86-10-62757497.The project was supported by the National Natural Science Foundation of China (50821061,20973003)and National Key Basic Research Special Foundation of China (2006CB806102).国家自然科学基金(50821061,20973003)和国家重点基础研究项目特别基金(2006CB806102)资助鬁Editorial office of Acta Physico -Chimica Sinica长波敏化发光铕配合物纳米粒子的制备与表征邵光胜薛富民韩荣成汤敏贤王远*(北京大学化学与分子工程学院,分子动态与稳态国家重点实验室,北京分子科学国家实验室,北京100871)摘要：以牛血清白蛋白(BSA)为保护剂,利用沉淀法制备了平均粒径为35nm 的Eu(tta)3dpbt (dpbt =2-(N ,N -二乙基苯胺-4-基)-4,6-二(3,5-二甲基吡唑-1-基)-1,3,5-三嗪,tta =噻吩甲酰三氟丙酮负离子)荧光纳米粒子.BSA 保护Eu(tta)3dpbt 纳米粒子在水中分散稳定性高,光稳定性好,长波敏化发光性能优良.其在可见光区激发峰位于415nm,激发峰尾部延展至470nm,发光量子产率为0.20(λex =415nm,25℃).在近红外双光子激发下可发出纯正的红光,Eu(tta)3dpbt 纳米粒子最大双光子激发作用截面为2.4×105GM (λex =830nm,1GM=10-50cm 4·s ·photo -1·particle -1).关键词：荧光纳米粒子;铕配合物;长波敏化发光;双光子激发;荧光生物探针中图分类号：O648Synthesis and Characterization of Europium Complex Nanoparticleswith Long -Wavelength Sensitized LuminescenceSHAO Guang -ShengXUE Fu -MinHAN Rong -ChengTANG Min -XianWANG Yuan *(Beijing National Laboratory for Molecular Sciences,State Key Laboratory for Structural Chemistry of Unstable and Stable Species,College of Chemistry and Molecular Engineering,Peking University,Beijing 100871,P.R.China )Abstract ：Luminescent nanoparticles (d av =35nm)of Eu (tta)3dpbt (dpbt =2-(N ,N -diethylanilin -4-yl)-4,6-bis (3,5-dimethylpyrazol -1-yl)-1,3,5-triazine,tta=thenoyltrifluoroacetonato)protected by bovine serum albumin (BSA)were prepared by a precipitation method.These particles disperse well in aqueous solutions to form stable and transparent colloidal solutions.The BSA -protected Eu(tta)3dpbt nanoparticles exhibit high photostability and excellent long -wavelength sensitized Eu III luminescence properties.The quantum yield of Eu III luminescence is 0.20(λex =415nm,25℃),and the excitation band for the Eu III luminescence centers at 415nm and extends to 470nm.Upon two -photon excitation,the prepared nanoparticles emit pure -red light with a two -photon excitation action cross section of 2.4×105GM (λex =830nm,1GM=10-50cm 4·s ·photo -1·particle -1).The prepared nanoparticles and the reported preparation method are promising for the development of luminescent bionanoprobes with excellent biocompatibility and luminescent properties.Key Words ：Luminescent nanoparticles;Europium complex;Long -wavelength sensitized luminescence;Two -photon excitation;Luminescent bioprobe稀土配合物独特的光学性质使其在生物标记[1-23]、非线性光学[24]、温度传感[25-27]、电致发光[28-30]等领域具有重要的应用价值.在生物标记方面,基于稀土发光的生物探针具有Stokes 位移大、发光寿命长和锐线发射等优点,据此发展起来的时间分辨荧光生化分析技术已被用于免疫分析[2-10]、DNA 检2031Acta Phys.-Chim.Sin.,2010Vol.26测[11-14]和生物成像[15-23].这项技术通过延时检测可有效地滤除背景信号的干扰,实现高灵敏度检测,同时避免了传统放射性标记方法带来的生物安全问题.由于稀土离子的f-f电子跃迁受跃迁选择定则的限制,稀土离子本身的光吸收能力很弱,要获得优良的光致发光性能,需要设计与稀土离子激发态能级匹配的天线配体,利用敏化作用(即由天线配体吸光并将能量传递至稀土离子)提高稀土离子的发光性能.许多稀土配合物在紫外光激发下具有优良的发光性质[31-34],然而,紫外激发光对生物样品损伤大,产生的背景荧光强,穿透深度浅,这使得紫外光激发发光的稀土荧光生物探针的应用受到较大限制.因此,将稀土配合物的激发窗口向长波方向扩展,研制具有优良可见光或近红外光激发发光性能的稀土配合物荧光探针成为稀土配合物设计、合成的重要方向[35-39].我们曾设计、合成了铕配合物Eu(tta)3dpbt (结构如图1所示)[35],Eu(tta)3dmbpt(dmbpt=2-(N,N-二乙基-2,6-二甲基苯胺-4-基)-4,6-二(3,5-二甲基吡唑-1-基)-1,3,5-三嗪)[37]以及Eu(tta)3bpt(bpt=2-(N,N-二乙基-4-基)-4,6-二(吡唑-1-基)-1,3,5-三嗪)[39]等一系列新型长波敏化铕离子发光配合物,它们具有优异的可见光及近红外双光子激发发光性能.时间分辨荧光光谱研究表明,Eu(tta)3dpbt的敏化传能主要通过单重态途径进行(即配体dpbt的三重态不参与传能),这是第一次以实验直接证实单重态途径可以成为稀土配合物的主要敏化途径.单重态能量传递机制摆脱了三重态能级的限制,使得Eu(tta)3dpbt类配合物具有很高的可见光敏化发光效率.最近,Wolfbeis等[25]利用Eu(tta)3dpbt等配合物发光强度的温敏性质研制了基于此类配合物的温度传感材料;我们发现与Eu(tta)3dpbt的温敏性质不同,Eu(tta)3bpt的发光量子产率在一定温度范围内随温度降低呈线性增加[39].Eu(tta)3dpbt等在可见光和近红外光激发下可高效发光的铕配合物为疏水性化合物,难以在水溶液中使用,这限制了它们在生物标记上的应用.解决此类问题的一个途径是制备能在水中稳定分散的所述配合物的荧光纳米粒子.用于制备稀土配合物纳米荧光生物探针的纳米粒子应具有以下特征:(1)在水溶液中具有良好的分散性和稳定性;(2)纳米粒子中的稀土配合物或其聚集体具有优良的光致发光性能,特别是在可见光及近红外光激发下的发光能力;(3)对激发光和发射光散射能力较弱;(4)具有较好的生物相容性.我们曾报道了分散于水-甲醇混合溶剂中,以十六烷基三甲基溴化铵(CTAB)为保护剂的Eu(tta)3dpbt 荧光纳米粒子[38],发现通过形成Eu(tta)3dpbt纳米粒子可以将Eu(tta)3dpbt激发窗口向长波方向拓宽,该纳米粒子在可见区的激发峰值位于420nm,较甲苯溶液中的Eu(tta)3dpbt的激发峰红移了18nm,激发带尾部延展到475nm,较甲苯溶液中的Eu(tta)3dpbt 的激发带向长波拓展了34nm,此性质为发展长波敏化发光稀土配合物材料提供了新的途径.上述纳米粒子具有优异的双光子激发发光性能,其最大双光子激发作用截面(δ·Φ)为3.2×105GM(λex=832 nm,1GM=10-50cm4·s·photo-1·particle-1).但是CTAB 的生物相容性不理想,而除去胶体溶液中的CTAB 则会导致纳米粒子分散稳定性下降,因此,研制生物相容性好、水溶液中分散稳定性高的Eu(tta)3dpbt荧光纳米粒子,对于发展纳米荧光生物探针具有积极的意义.牛血清白蛋白(BSA)是一种相对分子质量为68 kDa的球状蛋白,它含有580个氨基酸残基,是血液中疏水类维生素、激素、脂类化合物等物质的重要载体.牛血清白蛋白具有良好的生物相容性,不具有免疫原性,因而被广泛用作疏水性药物的载体[40].最近,BSA作为保护剂被用于合成具有良好生物相容性和较高荧光量子产率的金纳米粒子,此类金纳米粒子已被用于生物传感和细胞成像等领域[41-43].本文报道一种以BSA为保护剂的Eu(tta)3dpbt荧光纳米粒子及其制备方法,研究了此类纳米粒子的长波敏化铕离子发光性质.1实验部分1.1试剂图1Eu(tta)3dpbt结构式Fig.1Molecular structure of Eu(tta)3dpbt 2032No.7邵光胜等：长波敏化发光铕配合物纳米粒子的制备与表征Eu(tta)3dpbt 按照已报道方法合成[35];牛血清白蛋白购于Calbiochem 公司,丙酮(分析纯)购于北京试剂公司.实验用水采用18M Ω·cm 超纯水.1.2荧光纳米粒子的制备取2mL 浓度为1.6×10-4mol ·L -1的Eu(tta)3dpbt丙酮溶液,逐滴加入到8mL 浓度为0.5g ·L -1的BSA 水溶液中,将所得混合物在25℃下电磁搅拌10min,于25℃下旋转蒸发除去丙酮,以25000r ·min -1转速离心分离,将所得沉淀重新分散于8mL 含0.1%BSA 的10mmol ·L -1PBS 溶液中(磷酸盐缓冲溶液,pH =7.4),制得BSA 保护Eu(tta)3dpbt 荧光纳米粒子水溶胶.1.3表征纳米粒子的透射电子显微镜(TEM)照片采用日本Hitachi 公司生产的JEM 2000FX 透射电子显微镜拍摄.纳米粒子的水化半径和zeta 电位使用美国Brook heaven 公司生产的ZetaPlus Potential Analyz -ers 仪器测量.吸收光谱使用美国Varian 公司生产的CARY 1E 分光光度计记录.荧光光谱以及光稳定性测试使用日本Hitachi 公司生产的F -4500光谱仪记录,光稳定性测试使用150W 氙灯作为激发光源连续照射60min (光电倍增管电压400V ),记录荧光强度随照射时间的变化.荧光寿命使用英国Edinburgh Instruments 公司生产的FLS 920光谱仪测量.荧光量子产率(Φ)的测定按Demas 等[44]报道的方法进行,以DCM(4-dicyanomethylene -2-methyl -6-p -dimethylaminostyryl -4H -pyran)的正丙醇溶液作为参比(Φ=0.57±0.02).双光子激发作用截面测定以罗丹明B 作为参比,采用锁模Ti:Sapphire 激光器(Tsunami,日本Spectra Physics 公司生产)按相对双光子诱导荧光法[45]测定,双光子激发作用截面按如下公式计算(δ·Φ)S =F S F R ·C R nS S R!"·(δ·Φ)R其中,δ·Φ表示双光子激发作用截面,F 表示双光子诱导荧光强度,C 表示浓度,n 表示折光指数,下标S 和R 分别表示样品和参比物.2结果与讨论2.1荧光纳米粒子的形成牛血清白蛋白是良好的载脂蛋白,表面有多处由疏水残基构成的疏水区域,表面亲水基团使其在水中具有很大的溶解度.Eu(tta)3dpbt 为疏水性化合物,当将其丙酮溶液滴入BSA 水溶液中时,配合物分子迅速聚集形成Eu(tta)3dpbt 聚集体,BSA 表面的疏水残基与Eu(tta)3dpbt 聚集体之间的疏水相互作用使牛血清白蛋白吸附于Eu(tta)3dpbt 纳米粒子表面,而BSA 表面的亲水基团则使Eu(tta)3dpbt 纳米粒子稳定地分散于水溶液中,形成胶体溶液.本文制备的胶体溶液中Eu(tta)3dpbt 浓度为4.0×10-5mol ·L -1,将胶体溶液旋转蒸发除去丙酮后,所得水溶胶透明澄清,放置数月后无沉淀产生,表明BSA 对Eu(tta)3dpbt 纳米粒子具有很强的保护能力.图2为Eu(tta)3dpbt 纳米粒子的透射电子显微镜照片及其粒径分布图.所制备的荧光纳米粒子呈球状,平均粒径为35nm.动态光散射实验结果表明所制备的纳米粒子数均水化半径为50nm,比透射电镜法测定的平均粒径大15nm,表明吸附于纳米粒子表面的牛血清白蛋白的构象比较舒展,形成较厚的水化层.在PBS 缓冲溶液(pH=7.4)中,BSA 保护Eu(tta)3dpbt 纳米粒子和BSA 的zeta 电位分别为(-18.8±2.2)和(-14.1±1.1)mV,由此可见纳米粒子的荷电主要源于粒子表面吸附的BSA.由所制备纳米粒子的水化层和荷电情况可知,吸附于表面的BSA 导致的较大空间效应和纳米粒子的表面荷电是此类Eu(tta)3dpbt 纳米粒子在水中具有较好的分散稳定性的主要原因.2.2荧光纳米粒子的光学性质图3(a)中曲线1和2分别为BSA 保护Eu(tta)3dpbt纳米粒子胶体溶液和Eu(tta)3dpbt 甲苯溶液的紫外-可见吸收光谱.Eu(tta)3dpbt 分子聚集形成纳米粒子后,配体tta 的吸收峰位和峰形(约350nm)均未发生显著变化;甲苯溶液中Eu(tta)3dpbt 之dpbt 的吸收峰图2BSA 保护Eu(tta)3dpbt 纳米粒子的透射电子显微镜照片及其粒径分布图Fig.2TEM image and size distribution of the BSA -protected Eu(tta)3dpbt nanoparticles2033Acta Phys.-Chim.Sin.,2010Vol.26位于402nm,而BSA保护纳米粒子中Eu(tta)3dpbt 之dpbt的吸收峰位于415nm,即形成纳米粒子使Eu(tta)3dpbt在可见区的吸收峰红移了13nm.图3(a)中的曲线3和4分别为BSA保护Eu(tta)3dpbt纳米粒子与甲苯溶液中的Eu(tta)3dpbt分子的荧光激发光谱,纳米粒子中Eu(tta)3dpbt的荧光激发窗口延展至470nm,较甲苯溶液中Eu(tta)3dpbt的荧光激发带边红移约30nm.我们认为导致上述吸收峰红移现象的原因是由于在纳米粒子中形成了极性化合物Eu(tta)3dpbt的J-类聚集体.有机染料分子形成J-类聚集体的现象较为常见,J-类聚集体的形成往往会导致吸收峰和荧光激发峰红移.但是,稀土配合物形成J-类聚集体导致稀土荧光激发峰显著红移的现象却不多见[38].在CTAB保护Eu(tta)3dpbt纳米粒子中,Eu(tta)3dpbt 聚集体在可见区的吸收峰位于420nm;而在BSA 保护Eu(tta)3dpbt纳米粒子中,Eu(tta)3dpbt聚集体在可见区的吸收峰位于415nm.BSA保护Eu(tta)3dpbt 纳米粒子的荧光衰减动力学曲线如图3(b)所示(λex= 415nm,λem=614nm),对其进行双指数拟合(拟合相关系数0.999)给出两个荧光寿命,分别为0.33ms (56%)和0.10ms(44%).与此相对应,对CTAB保护Eu(tta)3dpbt纳米粒子荧光衰减动力学曲线进行拟合给出三个表观荧光寿命[38],分别为0.41、0.22和0.069ms.与纳米粒子的情况不同,甲苯溶液中Eu(tta)3dpbt分子的荧光衰减为单指数过程,其荧光寿命为0.48ms[35],这表明Eu(tta)3dpbt纳米粒子中的Eu(tta)3dpbt分子有多种堆积方式.BSA保护Eu(tta)3dpbt纳米粒子与CTAB保护Eu(tta)3dpbt纳米粒子在紫外-可见吸收光谱、荧光激发光谱及荧光寿命上的差异反映了两种纳米粒子中Eu(tta)3dpbt 分子的堆积方式有所不同.理想的纳米荧光生物探针对激发光和发射光的散射能力应该比较弱,最近Yuan等[10]利用三个含硅氧烷基团的化合物(以dpbt为天线配体的APS-BHHCT-Eu3+-dpbt,3-aminopropyl(triethoxyl)silane, tetraethyl orthosilicate)在反向微乳液中水解共缩合的方法合成了可被可见光激发发光的荧光纳米粒子并将其用于生物成像.此类纳米粒子对可见光(λ＞500nm)具有较强的散射能力,其紫外-可见吸收光谱中铕配合物的吸收峰较小,且叠加于强散射带之上.如图3(a)所示,BSA保护Eu(tta)3dpbt纳米粒子的紫外-可见吸收光谱表明其对可见光散射能力较弱.在25℃下,BSA保护Eu(tta)3dpbt纳米粒子的Eu III发光量子产率(Φ)为0.20(λex=415nm,未扣除纳米粒子散射信号的影响),415nm处的摩尔消光系数(ε)为3.3×104mol-1·L·cm-1,据此可计算出其发光亮度(ε·Φ)为6.6×103mol-1·L·cm-1,此值约为甲苯溶液中Eu(tta)3dpbt发光亮度的1/2(λex=415nm,25图3Eu(tta)3dpbt分子在BSA保护纳米粒子(实线)和甲苯溶液中(虚线)的光谱图Fig.3Spectra of BSA-protected Eu(tta)3dpbt molecules in the colloidal nanoparticles(solid line) and in toluene solution(dashed line)(a)UV-Vis absorption(lines1and2)and fluorescence excitation spectra(lines3and4)(λem=614nm);(b)decay curve of luminescence at614nm of the colloidal nanoparticles(λex=415nm),the solid line is the fit curve by the second order exponential decay;(c)photoluminescence spectra(λex=430nm).The Eu(tta)3dpbtconcentration in all of these samples is4.0×10-6mol·L-1. 2034No.7邵光胜等：长波敏化发光铕配合物纳米粒子的制备与表征℃).而在430nm光激发下,BSA保护Eu(tta)3dpbt纳米粒子的发光亮度为5.6×103mol-1·L·cm-1,纳米粒子中Eu(tta)3dpbt的发光强度为甲苯溶液中Eu(tta)3dpbt发光强度的2.8倍(图3(c)),显示出此类纳米粒子更为突出的长波敏化铕离子发光性质.光稳定性是纳米荧光生物探针的重要性质.如图4所示,BSA保护Eu(tta)3dpbt纳米粒子具有优良的光稳定性,经415nm激发光(150W氙灯)连续照射60min后,其发光强度(λem=614nm)仅下降至实验初始值的90%.1990年Webb等[46]将双光子激发荧光技术用于共聚焦显微成像,随后基于飞秒激光器的双光子激光扫描共聚焦显微镜也很快被投入使用.双光子荧光显微技术通常使用700-900nm波长的光源,一方面避开了紫外激发光对生物样品的损伤,另一方面由于激发光波长位于生物透光窗口,因而双光子显微技术穿透能力强、信噪比高,比单光子显微镜更适用于长时间观察和研究活体细胞和组织.具有优良双光子激发发光性能的铕配合物尚很有限, Eu(tta)3dpbt为迄今双光子敏化发光性质最好的铕配合物之一[36-39,47].BSA保护Eu(tta)3dpbt纳米粒子双光子激发谱峰位于830nm(图5),相对于甲苯溶液中的Eu(tta)3dpbt分子的双光子激发谱峰红移了22 nm.纳米粒子中Eu(tta)3dpbt分子的最大激发作用截面(δ·Φ)为24GM,远低于甲苯溶液中Eu(tta)3dpbt的最大激发作用截面(82GM).但一个纳米粒子却含有约1.0×104个配合物分子(Eu(tta)3dpbt的密度约为0.92g·cm-3[38],纳米粒子平均粒径为35nm),据此估算,其最大双光子激发作用截面约为2.4×105GM,该激发作用截面为见诸报道的CdSe/ZnS核壳量子点(d av≈4.5nm)[48]的5倍.3结论以BSA为保护剂制备了平均粒径为35nm的Eu(tta)3dpbt胶体粒子,吸附于Eu(tta)3dpbt粒子上的BSA使其表面荷负电并导致较强的空间稳定效应,此类纳米粒子在水溶液中具有良好的分散稳定性. BSA保护Eu(tta)3dpbt纳米粒子表现出优良长波敏化Eu III发光性质,其可见光激发峰值位于415nm,激发窗口延展至470nm,发光量子产率为0.20(λex= 415nm,25℃),双光子激发作用截面达2.4×105GM (λex=830nm,1GM=10-50cm4·s·photo-1·particle-1). BSA保护的Eu(tta)3dpbt纳米粒子与CTAB保护的Eu(tta)3dpbt纳米粒子在吸收峰位、荧光激发峰位和表观荧光寿命等方面具有显著区别,表明在两类纳米粒子中Eu(tta)3dpbt分子的堆积方式有较大差异.本文报道的荧光纳米粒子及其制备方法对发展生物相容性好、长波敏化发光性能优良的稀土配合物纳米荧光生物探针具有积极意义.References1Yuan,J.L.;Wang,G.L.Trends Anal.Chem.,2006,25:4902H覿rm覿,H.;Soukka,T.;L觟vgren,T.Clin.Chem.,2001,47:5613Soukka,T.;H覿rm覿,H.;Paukkunen,J.;L觟vgren,T.Anal.Chem., 2001,73:22544Soukka,T.;Paukkunen,J.;H覿rm覿,H.;L觟nnberg,T.;Lindroos,H.;L觟vgren,T.Clin.Chem.,2001,47:12695Blomberg,K.;Hurskainen,P.;Hemmil覿,I.Clin.Chem.,1999,45: 8556Tan,M.Q.;Wang,G.L.;Hai,X.D.;Ye,Z.Q.;Yuan,J.L.图5BSA保护Eu(tta)3dpbt纳米粒子的双光子激发作用截面(δ·Φ)和单光子荧光激发谱(λem=614nm)图(实线) Fig.5Two-photon excitation action cross sections (δ·Φ)and one-photon excitation spectrum (λem=614nm)(solid line)for the BSA-protectedEu(tta)3dpbt nanoparticlesThe experimental uncertainty for(δ·Φ)is about15%.1GM=10-50cm4·s·photo-1·particle-1图4BSA保护Eu(tta)3dpbt纳米粒子的光稳定性Fig.4Photostability curve of the BSA-protectedEu(tta)3dpbt nanoparticles2035Acta Phys.-Chim.Sin.,2010Vol.26J.Mater.Chem.,2004,14:28967Yuan,J.L.;Matsumoto,K.;Kimura,H.Anal.Chem.,1998,70: 5968Yuan,J.L.;Wang,G.L.;Kimura,H.;Matsumoto,K.Anal.Biochem.,1997,254:2839Wu,J.;Wang,G.L.;Jin,D.Y.;Yuan,J.L.;Guan,Y.F.;Piper,J.mun.,2008:36510Wu,J.;Ye,Z.Q.;Wang,G.L.;Jin,D.Y.;Yuan,J.L.;Guan,Y.F.;Piper,J.J.Mater.Chem.,2009,19:125811Chen,Y.;Chi,Y.M.;Wen,H.M.;Lu,Z.H.Anal.Chem.,2007, 79:96012Nurmi,J.;Wikman,T.;Karp,M.;L觟vgren,T.Anal.Chem.,2002, 74:352513Bortolin,S.;Christopoulos,T.K.;Verhaegen,M.Anal.Chem., 1996,68:83414Loannou,P.C.;Christopoulos,T.K.Anal.Chem.,1996,70:698 15Soini,E.;Pelliniemi,L.J.;Hemmil覿,I.;Mukkala,V.M.;Kankare, J.J.;Fr觟jdman,K.J.Histochem.Cytochem.,1988,36:144916Marriott,G.;Heidecker,M.;Diamandis,E.P.;Marriott,Y.Y.Biophys.J.,1994,67:95717Rulli,M.;Kuusisto,A.;Salo,J.;Kojola,H.;Simell,O.J.Immunol.Methods,1997,208:16918D′Aléo,A.;Pompidor,G.;Elena,B.;Vicat,J.;Baldeck,P.L.;Toupet.L.;Kahn,R.;Andraud,C.;Maury,O.ChemPhysChem,2007,8:212519Picot,A.;D′Aléo,A.;Baldeck,P.L.;Grishine,A.;Duperray,A.;Andraud,C.;Maury,O.J.Am.Chem.Soc.,2008,130:153220Law,G.L.;Wong,K.L.;Man,C.W.Y.;Wong,W.T.;Tsao,S.W.;Lam.M.H.W.;Lam.P.K.S.J.Am.Chem.Soc.,2008,130: 371421Kielar,F.;Congreve,A.;Law,G.L.;New,E.J.;Parker,D.;Wong, K.L.;Castreno,P.;De Mendoza,mun.,2008:2435 22Kielar,F.;Law,G.L.;New,E.J.;Parker,.Biomol.Chem., 2008,6:225623Law,G.L.;Wong,K.L.;Man,C.W.Y.;Tsao,S.W.;Wong,W.T.J.Biophoton.,2009,2:71824Andraud,C.;Maury,O.Eur.J.Inorg.Chem.,2009:435725Borisov,S.M.;Wolfbeis,O.S.Anal.Chem.,2006,78:509426Zelelow,B.;Khalil,G.E.;Phelan,G.;Carlson,B.;Gouterman,M.;Callis,J.B.;Dalton,L.R.Sens.Actuators B,2003,96:30427Stich,M.I.J.;Nagl,S.;Wolfbeis,O.S.;Henne,U.;Schaeferling,M.Adv.Funct.Mater.,2008,18:139928Vicinelli,V.;Ceroni,P.;Maestri,M.;Balzani,V.;Gorka,M.;V觟gtle,F.J.Am.Chem.Soc.,2002,124:646129Li,B.;Ma,D.G.;Zhang,H.J.;Zhao,X.J.;Ni,J.Z.Acta Phys.-Chim.Sin.,1998,14:305[李斌,马东阁,张洪杰,赵晓江,倪嘉缵.物理化学学报,1998,14:305]30Kido,J.;Okamoto,Y.Chem.Rev.,2002,102:235731Bunzli,J.C.G.;Piguet,C.Chem.Soc.Rev.,2005,34:104832Bunzli,J.C.G.;Comby,S.;Chauvin,A.S.;Vandevyver,C.D.B.J.Rare Earths,2007,25:25733He,G.S.;Tan,L.S.;Zheng,Q.;Prasad,P.N.Chem.Rev.,2008, 108:124534Ma,Y.;Wang,Y.Coord.Chem.Rev.,2010,254:97235Yang,C.;Fu,L.M.;Wang Y.;Zhang,J.P.;Wong,W.T.;Ai,X.C.;Qiao Y.F.;Zou,B.S.;Gui,L.L.Angew.Chem.Int.Edit.,2004,43:501036Fu,L.M.;Wen,X.F.;Ai,X.C.;Sun,Y.;Wu,Y.S.;Zhang,J.P.;Wang,Y.Angew.Chem.Int.Edit.,2005,44:74737Hao,R.;Li,M.Y.;Wang,Y.;Zhang,J.P.;Ma,Y.;Fu,L.M.;Wen,X.F.;Wu,Y.S.;Ai,X.C.;Zhang,S.;Wei,Y.G.Adv.Funct.Mater.,2007,17:366338Wen,X.F.;Li,M.Y.;Wang,Y.;Zhang,J.P.;Fu,L.M.;Hao,R.;Ma,Y.;Ai,ngmuir,2008,24:693239Xue,F.M.;Ma,Y.;Fu,L.M.;Hao,R.;Shao,G.S.;Tang,M.X.;Zhang,J.P.;Wang,Y.Phys.Chem.Chem.Phys.,2010,12:3195 40Patil,G.V.Drug Dev.Res.,2003,58:21941Tkachenko,A.G.;Xie,H.;Coleman,D.;Glomm,W.;Ryan,J.;Anderson,M.F.;Franzen,S.;Feldheim,D.L.J.Am.Chem.Soc,2003,125:470042Xie,J.;Lee,J.Y.;Wang,D.I.C.J.Phys.Chem.C,2007,111: 1022643Xie,J.P.;Zheng,Y.G.;Ying,J.Y.J.Am.Chem.Soc,2008,131: 88844Demas,J.N.;Crosby,G.A.J.Phys.Chem.,1971,75:99145Xu,C.;Webb,W.W.J.Opt.Soc.Am.B,1996,13:48146Denk,W.;Strickler,J.H.;Webb,W.W.Science,1990,248:73 47Werts,M.H.V.;Nerambourg,N.;Pélégry,D.;Grand,Y.L.;Blanchard-Desce,M.Photochem.Photobiol.Sci.,2005,4:53148Larson,D.R.;Zipfel,W.R.;Williams,R.M.;Clark,S.W.;Bruchez,M.P.;Wise,F.W.;Webb,W.W.Science,2003,300:14342036。