Starburst or Seyfert Using near-infrared spectroscopy to measure the activity in composite

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LINE FORMATION IN THE INNER STARBURST REGIONS OF AGN

W i n d s , B u b b l e s , & E x p l o s i o n s : A C o n f e r e n c e t o H o n o u r J o h n D y s o n . P át z c u a r o , M i c h o a c án , M éx i c o , 9-13 S e p t e m b e r 2002. E d i t o r s : S . J . A r t h u r & W . J . H e n n e y © C o p y r i g h t 2003: I n s t i t u t o d e A s t r o n o m ía , U n i v e r s i d a d N a c i o n a l A u t ón o m a d e M éx i c oRevMexAA (Serie de Conferencias),15,329–335(2003)LINE FORMATION IN THE INNER STARBURST REGIONS OF AGNI.AretxagaInstituto Nacional de Astrof´ısica,´Optica y Electr´o nica,M´e xicoRESUMENRevisamos la evidencia de que existan poblaciones estelares j´o venes en las regiones centrales (∼<200pc)de los n´u cleos gal´a cticos activos,y los mecanismos que los c´u mulos estelares tienen a su alcance para crear las l´ıneas de emisi´o n que caracterizan la actividad nuclear.ABSTRACTWe review the evidence for young stellar populations in the inner ∼<200pc of Active Galactic Nuclei (AGN),and the physical mechanisms through which the stars can potentially create the emission lines that characterize AGN.Key Words:ACTIVE GALACTIC NUCLEI —STARBURSTS —SUPERNOV AE1.INTRODUCTIONThe last two decades have seen a signiﬁcant ad-vance in our understanding of the phenomenology of various classes of AGN.The eﬀect of the nuclear orientation on the classiﬁcation of AGN has been particularly important (see Goodrich 2001for a re-view).Uniﬁed schemes postulate that the broad-line region (BLR)and the continuum emitting zone of an AGN are obscured,when observed edge-on,by a dusty obscuring torus with a size of order ∼1to 100pc.Under this scenario,and depending on the line of sight to the nucleus,identical objects can be classiﬁed as a Seyfert 2(Sy 2)or Seyfert 1(Sy 1).This conﬁguration can explain several observ-ables of Sy 2nuclei,for example the detection of broad lines in polarized light (Antonucci &Miller 1985;Miller &Goodrich 1990)which are directly detected in near-IR light (Goodrich,Veilleux,&Hill 1994),the presence of kpc-scale ionizing cones (Pogge 1988,1989;Tadhunter &Tsvetanov 1989),the large columns of neutral hydrogen that ab-sorb the X-ray emission (Mushotzky 1982;Bassani et al.1999;Risaliti,Miaulino,&Salvati 1999),and the cold far-IR colors (P´e rez-Garc´ıa,Rodr´ıguez-Espinosa,&Santolaya Rey 1998).However,the scheme,taken at face value,has some diﬃculty in in-corporating other observables:the high polarization levels observed in the broad polarized lines of Sy 2s compared to the modest polarization of the contin-uum (Goodrich &Miller 1989;Tran 1995;Cid Fer-nandes &Terlevich 1995),the similarity of the ob-served UV-continuum slopes of Sy 1and Sy 2s (Kin-ney et al.1991),the rich ∼100kpc environments sur-rounding Sy 2s but not Sy 1s (Dultzin-Hacyan et al.1999),the transient development of broad lines in otherwise classically quiescent type 2AGN (Storchi-Bergmann et al.1995;Aretxaga et al.1999a),and the absence of type 2QSOs (Hill,Goodrich,&De Poy 1995).Uniﬁed schemes,nevertheless,have survived 20years of detailed observational tests.Sensible varia-tions on the simple-minded statement “all Sy 2s are obscured Sy 1s seen edge-on”have been made,and the current consensus is that even if orientation is not a unique factor in creating the variety of all AGN types,it certainly is an important one to consider.2.YOUNG STELLAR POPULATIONS IN THEINNER REGIONS OF AGN The study of the ages of the stellar populations in the inner parts of AGN has been a source of debate over the last decade.We will deﬁne the inner regions as those at distances ∼<200pc from the gravitational center of the galaxy.In most cases,these regions are regarded as nuclear by the limitations of the reso-lution (∼1 )regularly achieved from ground-based optical facilities.Powerful circumnuclear starbursts at distances of ∼1kpc from the nucleus have been known to exist in many Seyfert galaxies for a long time (e.g.,NGC 1068).The nature of the stellar pop-ulations in the inner regions of the AGN,however,are still a matter of active research.In the early 90s the age determinations fo-cused on using near-IR absorption lines,where con-tamination by powerful line emission was minimal.The absorption features due to the Ca IR triplet λλ8494,8542,8662˚A in nearby Sy 2and Sy 1s were shown to be statistically stronger than the absorp-tions of normal spiral (S)and elliptical (E)galax-ies,once the possible presence of dilution by the power-law emission of an accretion disk had been accounted for (Terlevich,D´ıaz,&Terlevich 1990;329W i n d s , B u b b l e s , & E x p l o s i o n s : A C o n f e r e n c e t o H o n o u r J o h n D y s o n . P át z c u a r o , M i c h o a c án , M éx i c o , 9-13 S e p t e m b e r 2002. E d i t o r s : S . J . A r t h u r & W . J . H e n n e y © C o p y r i g h t 2003: I n s t i t u t o d e A s t r o n o m ía , U n i v e r s i d a d N a c i o n a l A u t ón o m a d e M éx i c o330ARETXAGAJim´e nez-Benito et al.2000;and see the ﬁgures by Nelson &Whittle in Terlevich 2001).This was inter-preted as direct evidence for a population of red su-pergiant stars that dominates the near-IR light.The high mass-to-light ratios L (1.6µm)/M ∼>3L /M inferred from the photospheric CO λ1.62,λ2.29µm absorptions of a sample of Sy 2nuclei conﬁrmed that red supergiants dominate the nuclear contin-uum emission in ∼50%of Sy 2s (Oliva et al.1995),whilst a similar conclusion could not be drawn for a Sy 1sample.With the advent of blue-sensitive detectors in many optical facilities,by the late 90s,the focus shifted to the detection of the UV and Balmer absorptions of massive stars.UV imaging and spectroscopy of 4Sy 2s selected by their strong [O III ]λ5007˚A and 1.4GHz ﬂuxes (which,in prin-ciple,are intrinsic AGN properties)show a ∼<100to 200pc resolved and broken-knot structure whose continuum spectrum is 100%that of a starburst,as derived from the strength of the absorption lines (Heckman et al.1997;Gonz´a lez Delgado et al.1998).The properties of the starbursts are L SB ≈1010to 1011L ,M SB ≈106to 107M and ages 3to 6Myr.Even in the low-luminosity type 2AGN (LINERs and Sy 2s)where the major component of UV light is unresolved,the continuum is totally dominated by the photospheres of OB stars (Maoz et al.1998;Col-ina et al.2002).The most complete surveys to date show that 50%out of 35Sy 2s have a continuum which is dominated by the emission of starbursts or post-starburst populations (Cid Fernandes et al.2001).A similar study for Sy 1s has not been at-tempted,due to the diﬃculty of decontaminating the UV-optical absorption-line spectrum from the corre-sponding emission-line spectrum.In smaller samples of radio-loud AGN a similar picture is starting to emerge:about 40%of the spec-tra of narrow-line radio galaxies show undiluted star-burst features (Aretxaga et al.2001;Tadhunter et al.2002),but the position of these starbursts cannot be determined to better than ∼1kpc from the center.Starbursts regions conﬁned to the inner ∼200pc can explain some of the paradoxes that the uniﬁed scheme faces,like the diﬀerent levels of polarization of broad lines and continuum and the similarity of the UV slopes between Sy 1and Sy 2s,if the AGN is not only obscured by the torus,but the torus is forming stars (Cid Fernandes &Terlevich 1995).A starburst-AGN connection has been proposed in at least three scenarios:starbursts giving birth to massive black holes (e.g.,Scoville &Norman 1988);black holes being fed by surrounding stellar clus-ters (e.g.,Perry &Dyson 1985;see also Pittard et al.2003);and also pure starbursts without black holes (e.g.,Terlevich &Melnick 1985;Terlevich et al.1992).The evidence for starbursts in Seyfert nuclei strongly supports some kind of connection.How-ever,it is still to be demonstrated that starbursts can explain some of the phenomenology characteris-tic of AGN,speciﬁcally the line emission spectrum,which ultimately is what deﬁnes an AGN (Seyfert 1943;Baldwin,Phillips,&Terlevich 1981).3.AGN NARROW-LINE EMISSION SPECTRUMFROM STARBURSTS Early attempts to reproduce the AGN emission-line spectrum,using stars as the only source of ion-ization,have met with only moderate success and much controversy.Terlevich &Melnick (1985)pro-posed the existence of evolved massive stars of T eﬀ∼>100,000K (extremely hot WC or WO stars,which they named warmers ),by directly applying the stel-lar evolutionary models in vogue at the time.These stars,present in a 3to 8Myr starburst,had enough hard-UV photons to reproduce the diagnostic lines of AGN.It was the addition of an optically thick stel-lar wind to the stellar interior models (Maeder 1990;Schaerer et al.1993)that suggested the extremely-luminous blue phase did not exist,and since then the idea of warmers as a source of ionization in AGN has slowly been disregarded by its proposers (Cid Fernandes 1997;Terlevich 2001).Normal starbursts can reproduce some of the lines observed in AGN.In particular,the inferred properties of the nuclear/circumnuclear starbursts that reproduce the observed UV absorption spec-tra of Sy 2s (§2)have enough photons to ionize the whole Balmer emission series (Gonz´a lez Delgado et al.1998;Colina et al.2002),but the photospheres of those young stars cannot,in general,account for the higher-ionization emission-line species,like [Ne V ]λ3426˚A or He II λ4686.Low-ionization species,however,can be repro-duced with the ionizing power of hot stars.A star-burst with normal OB stars,in dense media (∼>5×103cm −3),can reproduce weak [O I ]λ6300˚A /H αLINERs (Shields 1992;Filippenko &Terlevich 1992).4.TYPE IIN SNE OR COMPACT SNREMNANTS The modest success of using only stars to ion-ize the gas in an AGN was soon to be replaced by the by-products of their evolution in pure starburst models:the supernova explosions and the quick re-processing of their kinetic energy by dense circum-W i n d s , B u b b l e s , & E x p l o s i o n s : A C o n f e r e n c e t o H o n o u r J o h n D y s o n . P át z c u a r o , M i c h o a c án , M éx i c o , 9-13 S e p t e m b e r 2002. E d i t o r s : S . J . A r t h u r & W . J . H e n n e y © C o p y r i g h t 2003: I n s t i t u t o d e A s t r o n o m ía , U n i v e r s i d a d N a c i o n a l A u t ón o m a d e M éx i c oLINE FORMATION IN THE INNER STARBURST REGIONS OF AGN331Fig. 1.Rest-frame optical spectra of a collection of type IIn SNe (Di Carlo et al.2002):constant +log ﬂux (erg s −1cm −2˚A −1)versus wavelength (˚A ).stellar media.These particular SNe were ﬁrst con-sidered by Terlevich,Melnick,&Moles (1987)when they extrapolated the models of remnants evolving in dense media of ∼105cm −3(Shull 1980;Wheeler,Mazurek,&Sivaramakrishnan 1980).SNe exploding in much denser media were soon found in the outer regions of nearby spiral galaxies (Filippenko 1989;Stathakis &Sadler 1991)and gave rise to a diﬀer-ent SN spectroscopic class:type IIn SNe (Schlegel 1990).These objects can potentially explain the high-ionization narrow emission lines seen in type 2AGN (§6),and also provide the broad emission lines and UV–optical–IR light variations,which,over long timescales (∼>1month),resemble those of type 1AGN (§5).The spectra of SN IIn are characterized by the presence of prominent narrow emission lines (hence the “n”)sitting on top of broad components withFWHM ∼<15,000km s−1at maximum light (see Figure 1),and look extremely similar to the spec-tra of type 1AGN (see a comparison in Filip-penko 1989and Terlevich 2001).They do not show the characteristic broad P Cygni signatures of stan-dard SNe,although narrow P Cygni proﬁles are de-tected in some cases at high spectral resolution (e.g.,SN 1997ab:Salamanca et al.1998).SN IIn are nor-mally associated with regions of recent star forma-tion (Schlegel 1990).Despite these general charac-teristics,SN IIn as a group exhibit considerable het-erogeneity (see also Filippenko 1997):•Some type IIn SNe have an extremely slow de-cay of luminosity after maximum light,which makes them,after 600days,approximately 5mag brighter in the V -band than standard SN IIP or SN IIL (e.g.,SN 1988Z:Stathakis &Sadler 1991,see Figure 2);however,others have a photometric behaviour much like standard type IIL SNe (e.g.,SN 1999el:Di Carlo et al.2002).•Their peak luminosities (M V ∼−18.8)are within the range of classical type IIP SNe(Richard-Fig.2.V -band light curves of 16type IIn SNe with known optical maxima,normalized to maximum light (Aretxaga et al.2003).son et al.2002),and thus they are not particularly overluminous at optical wavelengths.•Some of those SNe that decay slowly are probably hypernovae,with kinetic energies in the range of ∼1052erg (SN 1988Z:Aretxaga et al.1999b;SN 1997cy:Turatto et al.2000;SN 1999E:Rigon et al.2003),while other slowly decaying SNe have modest integrated energies of ∼1049erg (e.g.,SN 1995N:Pastorello et al.2003).•Among the energetic type IIn SNe,two are probably associated with gamma-ray bursts (SN 1997cy:Germany et al.2000;SN 1999E:Rigon et al.2003);•Extremely bright radio and X-ray emission has been detected in some type IIn SNe (e.g.,SN 1988Z:van Dyk et al.1993,Fabian &Terlevich 1996;SN 1995N:Fox et al.2000),but emission at these wavelengths is not common in others (e.g.,SN 1997ab).•Whenever the forbidden-line ratios have been used to estimate the density of the narrow-line pro-ducing region in type II SNe,values in the range 106to 109cm −3have been found (e.g.,SN 1988Z:Stathakis &Sadler 1991;SN 1995N:Fransson et al.2002.;SN 1995G:Pastorello et al.2002).It was soon recognized that the special charac-teristics of type IIn SNe are due to the strong inter-action of the ejecta from the explosion with a dense circumstellar medium (Chugai 1991;Terlevich et al.1992),the origin of which is probably the compressed winds of the progenitor star:slow decays and broadvariable lines originate in the dense (∼>1012cm −3)double shell structure created by the outer and in-ner shocks as they sweep the dense (∼107cm −3)circumstellar medium and the ejecta,respectively;and the narrow lines are produced by the unshockedW i n d s , B u b b l e s , & E x p l o s i o n s : A C o n f e r e n c e t o H o n o u r J o h n D y s o n . P át z c u a r o , M i c h o a c án , M éx i c o , 9-13 S e p t e m b e r 2002. E d i t o r s : S . J . A r t h u r & W . J . H e n n e y © C o p y r i g h t 2003: I n s t i t u t o d e A s t r o n o m ía , U n i v e r s i d a d N a c i o n a l A u t ón o m a d e M éx i c o332ARETXAGAcircumstellar medium,which is ionized by the radi-ation coming from the shocks.In the case of strong interactions,like in SN 1988Z,where the estimated radiated energy from the radio to X-rays in the ﬁrst 10years of evolution exceeds 3×1051erg,and is probably close to 1052erg (Aretxaga et al.1999b),i.e.,two orders of magni-tude larger than typical SN events,the name “su-pernova”does not do justice to the phenomenology we witness.The large radiated energies imply that most of the kinetic energy released in the explosion must be reprocessed into radiation within the ﬁrst decade of evolution,much as classical SN remnants behave over the course of thousands of years.These type IIn SNe are referred to as “compact supernova remnants”(cSNRs)by Terlevich et al.(1992),which describe events where the energetics are dominated by the conversion of kinetic into radiated energy,and not by the thermal cooling of the expanding atmo-sphere of an exploding star.5.TYPE IIN SNE IN THE INNERCIRCUMNUCLEAR REGIONS OF AGN?There is little doubt that if a SN IIn explodes in the center of a normal galaxy,the nucleus would be classiﬁed as a Sy 1while the prominent broad lines remain visible.In fact,there has been a suc-cession of theoretical studies that attempt to explain the phenomenology of lines and continuum at UV to optical wavelengths in Seyfert 1nuclei in terms of a starburst that undergoes SN IIn explosions (Ter-levich et al.1987,1992,1995;Aretxaga &Terlevich 1994;Aretxaga,Cid Fernandes,&Terlevich 1997).Massive starbursts in the center of ∼50%of type 2AGN have been discovered (§2).However,it is still to be determined whether they also populate the centers of type 1AGN.The intensity of the cal-cium triplet absorptions provides some evidence in this direction (Jim´e nez-Benito et al.2000).If starbursts are indeed present in the nuclei of type 1AGN,and have similar ages to those of type 2AGN,they will produce a considerable number of SNe,and if these SNe are type IIn,then they po-tentially can reproduce the type 1AGN phenomenol-ogy at IR–optical–UV wavelengths.Since starbursts are subject to a scaling relationship,where the SN rate (νSN )and the optical luminosity coming from stars (L ∗B )are related along the lifetime of the SN II explosion phase (∼10to 70Myr)by νSN /L ∗B ≈2×10−11yr −1L −1B (Aretxaga &Terlevich 1994),the rates required to reproduce standard Sy 1s with pure starbursts are between 0.2SN yr −1for AGN of lumi-nosities close to NGC 4151and 0.5SN yr −1for thoseclose to NGC 5548.These correspond to masses of the starbursts around 108to 109M ,depending on the slope and lower end of the initial mass function.These proposed starbursts are thus 1to 2orders of magnitude more massive than those found in Sy 2s to date.The light curve that a cluster of luminosity similar to the nucleus of NGC 4151would produce,if all SNe were type-IIn,is represented in Figure 3together with the historically observed light curve of NGC 4151for comparison.The simulation,al-though not identical to the observed object (since it is produced stochastically),can reproduce the basic long-term properties such as mean luminosity,rms and power spectrum,with just one free parameter:the density of the circumstellar medium in which the cSNRs evolve (Aretxaga &Terlevich 1994).The theoretical models of cSNRs show that the UV and Balmer lines respond to variations of the continuum on timescales of a few days to several tens of days,depending on the line species,and that these lags are similar to those found in NGC 5548,but they are not created by time-travel delays (Terlevich et al.1995).However,the lags have not been explored in real type IIn SNe since the sampling of the light vari-ations in these sources is still very scarce.The only lag clearly detected is that of the H αline in SN 1988Z (Turatto et al.1993),but the lag is ∼200days,and these values are not seen (although they could have been detected)in other type IIn.The case of NGC 7582,a classical Sy 2that sud-denly mutated into a Sy 1(Aretxaga et al.1999a),is probably one of the most compelling examples where the SN explanation works,although this is by no means a unique solution.Many of the nuclear prop-erties of NGC 7582support a uniﬁed scheme where the true Sy 1nature is hidden by an obscuring torus:a sharp-edged [O III ]outﬂow in the form of a cone is observed (Morris et al.1985);optical spectropo-larimetry does not reveal a hidden broad-line re-gion,but since the far-IR colors 60µm to 25µm are very red,the absence has been taken as support for an edge-on thick torus able to block even the light scattered towards the observer (Heisler,Lumsden,&Bailey 1997);indeed,a large column density of neu-tral H (N H ∼1024cm −2)also blocks the hard X-rays (Warwick et al.1993),but this absorption is variable and decreased (∆N H ∼1023cm −2)at the time of the transition between Seyfert types (Turner et al.2000).The presence of stars in the nucleus is also ﬁrmly es-tablished:Morris et al.(1985)found a steep gradient of H αperpendicular to the [O III ]cone,which they interpret as a 1kpc disk of H II regions oriented at 60◦from the plane of the galaxy;the CO absorptionW i n d s , B u b b l e s , & E x p l o s i o n s : A C o n f e r e n c e t o H o n o u r J o h n D y s o n . P át z c u a r o , M i c h o a c án , M éx i c o , 9-13 S e p t e m b e r 2002. E d i t o r s : S . J . A r t h u r & W . J . H e n n e y © C o p y r i g h t 2003: I n s t i t u t o d e A s t r o n o m ía , U n i v e r s i d a d N a c i o n a l A u t ón o m a d e M éx i c oLINE FORMATION IN THE INNER STARBURST REGIONS OF AGN333(a )(b )Fig.3.(a )Light curve of NGC 4151in absolute B -band magnitudes.The uncertainties in the derivation of the radial velocity of the galaxy are shown as the point with error bars at the upper-right corner of the panel.(b )Simulated light curve for a starburst which undergoes a SN rate νSN =0.3yr −1,with energy of explosions E =(3.0±0.1)×1051erg evolving in a medium of n =(1.0±0.3)×107cm −3.The sampling of points in the simulated light curve is identical to that in the top panel.Observational errors of σobs =0.1mag are included.The dotted lines represent the simulated light curve (Aretxaga &Terlevich 1994).lines and large near-IR light-to-mass ratio are simi-lar to those of H II galaxies and a factor of 5larger than those of normal galaxies,indicating that red supergiants dominate the light of the inner 200pc at those wavelengths (Oliva et al.1995);the expected SN rate of this starburst is 0.02SN yr −1(Aretxaga et al.1999a).The light variations in the optical were not consistent with a standard reddening law vari-ation (the Goodrich 1989test),which could be the result of nuclear light traveling through a region of less obscuration in the torus.The ﬂare,which was followed for a few months,had a similar decaying law as that of SN 1983K (retrospectively classiﬁed as a type IIn,and a fast decayer like SN 1999el)and line-width variations of ∼12,000km s −1to 5000km s −1in a few months are also typical of type IIn (Aretx-aga et al.1999a).The X-ray data taken during the ﬂare are consistent with both explanations,either a change in reddening in the torus or a type IIn SN onset (Turner et al.2000).6.DISCUSSIONStarbursts in type 2AGN have been directly found on scales which could correspond to the sur-roundings of a dusty torus.We also could be seeing the outskirts of a massive cluster embedded within the dusty torus.The stars in the best studied cases (e.g.,Gonz´a lez Delgado et al.1998)can reproduce all the observed Balmer line emission and contin-uum,but they do not have enough ionizing photons to produce the high-ionization lines that character-ize AGN.Some type IIn SNe provide hard energy photons in suﬃcient quantities to produce the high-ionization lines of AGN,and also coronal lines like [Fe X ]λ6375˚A ,[Fe VII ]λ5159˚A ,[Fe VII ]λ6086˚A ...(Turatto et al.1993;Fox et al.2000).If these SNe are mixed with the dust,or if they are located in the central region of the AGN and are hidden by the torus,they could reproduce the rest of the ionization characteristics of the Sy 2s.The black hole should still be responsible for properties such as the rela-tivistic broad Fe K αlines that are present in many Sy 2s (Turner et al.1997),but it does not necessarily have to be responsible for the UV-optical emission.In turn,some type IIn SNe in the outskirts of the torus could be directly visible,mimicking the phenomenology of AGN (e.g.,NGC 7582).If one is to extrapolate this result to classical Sy 1,the masses of the stellar clusters required are 1to 2orders of magnitude larger than those di-rectly detected in type 2AGN.These could per se be larger,or we could be seeing more of the cluster as part or all of the dust in Sy 1s is blown up or dis-persed,as suggested by Dultzin-Hacyan et al.(1999).Type IIn SNe do mimic many of the characteristics of AGN,but there are still many of their properties that are unexplored,e.g.,the lags of their emission lines or the small-scale variations.The lag ∝L 1/2relationship found in AGN is one of the most out-standing predictions and successes of the standard model of AGN (e.g.,Wandel,Peterson,&Malkan 1999),where the lines originate in (or near)an ac-cretion disk that surrounds the supermassive black hole.Observationally little is known about lags in type IIn SNe.Only SN 1988Z shows a delay of theW i n d s , B u b b l e s , & E x p l o s i o n s : A C o n f e r e n c e t o H o n o u r J o h n D y s o n . P át z c u a r o , M i c h o a c án , M éx i c o , 9-13 S e p t e m b e r 2002. E d i t o r s : S . J . A r t h u r & W . J . H e n n e y© C o p y r i g h t 2003: I n s t i t u t o d e A s t r o n o m ía , U n i v e r s i d a d N a c i o n a l A u t ón o m a d e M éx i c o334ARETXAGAH αresponse to the continuum,of ∼200days.In the rest of the SNe these large lags are not seen,and smaller ones,if present,cannot be character-ized with the available data.A variation of the lag with the environment of the circumstellar material is predicted,and one can conceive of a link with the to-tal luminosity of the cluster.The details,however,have not been worked out,and most importantly,they haven’t been tested against bona-ﬁde isolated SN IIn.It might also be the case that the stellar clusters in AGN do not produce type IIn SNe at all,but in-stead normal type IIs.It would still be important,however,to learn how to distinguish the optical vari-ations of type IIn SNe from those of accretion pro-cesses in AGN,and estimate the true SN rate in the clusters.IA’s research is partly supported by CONACyT grant E-32143.REFERENCESAntonucci,R.R.J.,&Miller,J.S.1985,ApJ,297,621Aretxaga,I.,Cid Fernandes,R.,&Terlevich,R.J.1997,MNRAS,286,271Aretxaga,I.,Joguet,B.,Kunth,D.,Melnick,J.,&Ter-levich,R.J.1999a,ApJL,519,123Aretxaga,I.,&Terlevich,R.J.1994,MNRAS,269,462Aretxaga,I.,Terlevich,E.,Terlevich,R.J.,Cotter,G.,&D´ıaz,A.I.2001,MNRAS,325,636Aretxaga,I.,et al.1999b,MNRAS,309,343.2003,in preparationBaldwin,J.A.,Phillips,M.M.,&Terlevich,R.,1981,PASP,93,5Bassani,L.,et al.,1999,ApJS,121,473Chugai,N.N.1991,MNRAS,250,513Cid Fernandes,R.1997,RevMexAA(SC),6,201Cid Fernandes,R.,Heckman,T.,Schmitt,H.,Gonz´a lezDelgado,R.M.,&Storchi-Bergmann,T.2001,ApJ,558,81Cid Fernandes,R.,&Terlevich,R.J.1995,MNRAS,272,423Colina,L.,Gonz´a lez-Delgado,R.M.,Mas-Hesse,J.M.,&Leitherer,C.2002,ApJ,579,545Di Carlo,E.,et al.2002,ApJ,573,144Dultzin-Hacyan,D.,Krongold,Y.,Fuentes-Guridi,I.,&Marziani,P.1999,ApJ,513,111Fabian,A.C.,&Terlevich,R.J.,1996,MNRAS,280,L5Filippenko,A.V.1989,AJ,97,726.1997,ARA&A,35,309Filippenko,A.V.,&Terlevich,R.J.1992,ApJ,397,L79Fox,D.W.,et al.2000,MNRAS,319,1154Fransson,C.,et al.2002,ApJ,572,350Germany,L.M.,Reiss,D.J.,Sadler,E.M.,Schmidt,B.P.,&Stubbs,C.W.2000,ApJ,533,320Gonz´a lez Delgado,R.M.,et al.1998,ApJ,505,174Goodrich,R.W.1989,ApJ,340,190.2001,in The Starburst-AGN Connection,eds.I.Aretxaga,D.Kunth,&R.M´u jica (Singapore:World Scientiﬁc),69Goodrich,R.W.,&Miller,J.S.1989,ApJ,346,21Goodrich,R.W.,Veilleux,S.,&Hill,G.J.1994,ApJ,422,521Heckman,T.M.,et al.1997,ApJ,482,144Heisler,C.A.,Lumsden,S.L.,&Bailey,J.A.1997,Nature,385,700Hill,G.J.,Goodrich,R.W.,&De Poy,D.L.1995,ApJ,462,163Jim´e nez-Benito,L.,D´ıaz,A.I.,Terlevich,R.,&Ter-levich,E.2000,MNRAS,317,907Kinney,A.L.,Antonucci,R.R.J.,Ward,M.J.,Wilson,A.S.,&Whittle,M.1991,ApJ,377,100Maeder,A.1990,A&AS,84,139Maoz,D.,Koratkar,A.,Shields,J.C.,Ho,L.C.,Filip-penko,A.V.,&Sternberg,A.1998,AJ,116,55Miller,J.S.,&Goodrich,R.W.1990,ApJ,355,456Morris,S.,Ward,M.,Whittle,M.,Wilson,A.S.,&Tay-lor,K.1985,MNRAS,216,193Mushotzky,R.F.1982,ApJ,256,92Oliva,E.,Origlia,L.,Kotilainen,J.K.,&Moorwood,A.F.M.1995,A&A,301,55P´e rez-Garc´ıa,A.M.,Rodr´ıguez-Espinosa,J.M.,&San-tolaya Rey,A.E.1998,ApJ,500,685Perry,J.,&Dyson,J.E.1985,MNRAS,213,665Pastorello,A.,et al.2002,MNRAS,333,27.2003,in preparationPittard,J.M.,Dyson,J. E.,Falle,S. A. E.G.,&Hartquist,T.W.2003,RevMexAA(SC),15,317(this volume)Pogge,R.W.1988,ApJ,328,519.1989,ApJ,345,730Richardson,D.,Branch,D.,Casebeer,D.,Millard,J.,Thomas,R.C.,&Baron,E.2002,AJ,123,745Rigon,L.,et al.2003,MNRAS,in pressRisaliti,G.,Maiolino R.,&Salvati M.1999,ApJ,522,157Salamanca,I.,Cid Fernandes,R.,Tenorio-Tagle,G.,Telles,E.,Terlevich,R.J.,&Mu˜n oz-Tu˜n ´o n,C.1998,MNRAS,300,L17Schaerer,D.,Meynet,G.,Maeder,A.,&Schaller,G.1993,A&AS,98,523Schlegel,E.M.1990,MNRAS,244,269Scoville,N.,&Norman,C.1988,ApJ,332,163Seyfert,C.K.1943,ApJ,97,28Shields,J.C.1992,ApJ,399,L27Shull,J.M.1980,ApJ,237,769Stathakis,R.A.,&Sadler,E.M.1991,MNRAS,250,786Storchi-Bergmann,T.,Eracleous,M.,Livio,M.,Wilson,A.S.,Filippenko,A.V.,&Halpern,J.P.1995,ApJ,443,617Tadhunter,C.,&Tsvetanov,Z.1989,Nature,341,422Tadhunter,C.,et al.2002,MNRAS,330,977Terlevich,E.,D´ıaz,A.I.,&Terlevich,R.J.1990,MN-。

高三英语天文观测设备单选题50题

高三英语天文观测设备单选题50题1. An ______ is a building or place equipped with telescopes and other instruments for observing astronomical objects.A. observatoryB. laboratoryC. factoryD. library答案：A。

解析：本题考查名词词义辨析。

observatory意为天文台，是配备望远镜等仪器用于观测天文物体的建筑或场所，符合题意。

laboratory是实验室，主要用于科学实验；factory是工厂，用于生产制造；library是图书馆，用于藏书和供人阅读学习，这三个选项均不符合天文观测场景的描述。

2. The ______ is an important tool for astronomers to observe the stars and galaxies far away.A. microscopeB. telescopeC. magnifierD. binoculars答案：B。

解析：本题考查天文观测工具相关的名词。

telescope望远镜是天文学家观测遥远恒星和星系的重要工具。

microscope是显微镜，用于观察微小的物体，如细胞等；magnifier是放大镜，主要用于放大近距离的小物体；binoculars是双筒望远镜，虽然也可用于观测，但在天文观测中telescope更为专业和常用。

3. In the observatory, the ______ of the telescope needs to be adjusted precisely to get a clear view of the celestial bodies.A. lensB. buttonC. handleD. box答案：A。

解析：本题考查名词在天文观测设备中的部件。

The scatter in the near-infrared colour-magnitude relation in spiral galaxies

a r X i v :a s t r o -p h /9808232v 1 21 A u g 1998Mon.Not.R.Astron.Soc.000,000–000(0000)Printed 1February 2008(MN L A T E X style ﬁle v1.4)The scatter in the near-infrared colour-magnitude relationin spiral galaxiesR.F.Peletier 1,2and R.de Grijs 3,21Dept.of Physics,University of Durham,South Road,Durham DH13LE,UK2Kapteyn Astronomical Institute,Postbus 800,9700AV Groningen,The Netherlands3Astronomy Department,University of Virginia,PO Box 3818,Charlottesville,VA 22903-0818,USA1February 2008Key words:galaxies:fundamental parameters;galaxies:photometry;galaxies:struc-tureABSTRACTWe have determined a dust-free colour-magnitude (CM)relation for spiral galaxies,by using I −K colours in edge-on galaxies above the plane.We ﬁnd that the scatter in this relation is small and approximately as large as can be explained by observational uncertainties.The slope of the near-IR CM relation is steeper for spirals than for elliptical galaxies.We suggest two possible explanations.First,the diﬀerence could be caused by vertical colour gradients in spiral galaxies.In that case these gradients should be similar for all galaxies,on average ∼0.15dex in [Fe/H]per scale height,and increase for later galaxy types.The most likely explanation,however,is that spirals and ellipticals have intrinsically diﬀerent CM relations.This means that the stars in spirals are younger than those in ellipticals.The age,however,or the fraction of young stars in spiral galaxies,would be determined solely by the galaxy’s luminosity,and not by its environment.1INTRODUCTIONBaum (1959)and de Vaucouleurs (1961)ﬁrst established that early-type galaxies obey a tightly constrained colour –absolute magnitude (CM)relation.Somewhat later,in 1978,Sandage &Visvanathan ﬁrst proposed that the relation,if applied to galaxy clusters,could provide a signiﬁcant con-straint on their past history of star ing mod-ern detectors Bower et al.(1992a,b)determined that both in Virgo and Coma the scatter around the CM relations in U −V ,B −V and V −K is extremely small (≤0.05mag),comparable to their observational uncertainties (∼0.03mag).This allowed them to put strong constraints on the age scatter of early-type galaxies in clusters.Recently,Ellis et al.(1997)again found very small scatter in the CM diagram of early-type galaxies in clusters at z ∼0.54,which led them to conclude that most of the star formation in these clusters must have ended at z ∼3.These and other obser-vations have led people to attribute the origin of the CM relation to changes in metallicity,caused by the fact that the larger galaxies have larger binding energies,so that the gas can be enriched to higher metallicities (e.g.,Mathews &Baker,1971;Faber,1977;Arimoto &Yoshii,1987).For spiral galaxies,however,the situation is much more complicated,because of recent star formation,as well as ex-tinction by dust.Visvanathan &Griersmith (1977)found forearly-type spiral galaxies in the Virgo cluster (S0/a to Sab)within the errors the same optical CM relation as had been found for E/S0galaxies,except that the scatter was ter Tully et al.(1982)and Mobasher et al.(1986)estab-lished optical-IR and IR CM relations for early and late type spirals,with considerable scatter.Tully et al.(1982)claimed that the diﬀerence between early and late type spirals could be explained by the presence of more star formation in late type spirals,aﬀecting especially the blue light,and hence the B −H colour.In the last decade this picture has not changed very much.Interpreting a CM relation using the B −H colour is extremely complicated,given the fact that the eﬀects of star formation and extinction counteract each other.Valentijn (1990)and afterwards many others authors have shown that the extinction corrections applied by Tully et al.are a factor of 2or 3too low,but in any case very un-certain.To be able to draw more detailed conclusions about galaxy formation and evolution,a diﬀerent way of present-ing the CM relation has to be found.In this paper we argue that this can be done using the I −K colour,obtained at po-sitions in the galaxy that are likely to be very little aﬀected by dust extinction.The main diﬀerence with previous work is that we use local colours here,which is now possible due to advances in IR detector technology.Integrated colours from,e.g.,aperture photometry are very hard to interpret due to the complicating eﬀects of extinction and stellar populationc0000RAS2R.F.Peletier and R.de Grijsdiﬀerences,except if the colours do not change much inside the galaxy,as is the case in elliptical galaxies.Secondly,to eliminate the eﬀects of very recent star formation as much as possible,we have tried to use a colour that is as red as possible,while trying to keep the wavelength baseline large. Knapen et al.(1995)show convincingly that the I−K colour is predominantly a dust/extinction indicator,while recent star formation shows up much less than in a blue colour like B−V.In this paper we try to use the I−K vs.M K rela-tion to determine a colour-magnitude relation with as low a scatter as possible for spiral galaxies.This relation can be used to study the following problems:1)The nature of spiral galaxies themselves.By comparing spirals with ellip-ticals we can study the star formation history of spirals in a very direct way,making only the simple assumption that elliptical galaxies are coeval(see,e.g.,Kodama et al.1998).2)The role of the environment on the evolution of galaxies.A tight CM relation for spirals can serve as a very useful tool to study the evolution of spiral galaxies and the role of their environment and3)The use of the CM diagram as a distance indicator,very useful for isolated galaxies in the ﬁeld(de Grijs&Peletier,in preparation).In the following section we present the data used for our CM relation.In section3the scatter on the relation is dis-cussed,and a comparison is made with early-type galaxies. In section4we discuss some implications of the low scatter for the nature of spiral galaxies,after which we summarise the paper.2A DUST-FREE CM RELATIONTo determine a dust-free CM diagram we decided to limit ourselves to two datasets,of which we can be fairly conﬁdent that the amount of extinction by dust is very small.Theﬁrst is a sample of22edge-on galaxies,described in de Grijs et al.(1997).This is a random subsample of Southern non-interacting galaxies with inclinations larger than87o and blue diameters D B25larger than2’.2.These galaxies were ob-served in I and K.More observational details can be found in de Grijs(1998).It was found that although the central regions are heavily obscured,the vertical colour proﬁles are symmetric,featureless and with colour almost constant as a function of radius(de Grijs et al.1997).These3properties together indicate that the colours here are not reddened by dust.We therefore took the average of the colour on both sides of the galaxy in the region where the colour proﬁle is featureless.Absolute K-magnitudes were determined using a simple Hubbleﬂow model correcting for the motion of the local group(see de Grijs1997).The second sample is the sample of early-type spiral galaxies of Peletier&Balcells (1997).This is a sample of galaxies with inclinations larger than50o,for which the colours were determined in the bulge opposite to the dustlane.The fact that the colour proﬁles on one side of the galaxy generally are featureless shows that our assumption about negligible extinction is justiﬁed.How-ever,since for some of the latest-type galaxies of this sample the bulges are small,and therefore their colours probably more aﬀected by extinction,we only consider the earliest-type spirals(type≤1)for the analysis done in this paper. Other samples are available(e.g.,de Jong1996),but less is known about their extinction properties,and by including them the scatter generally increases.The colours of both samples used were corrected for Galactic extinction and redshift in the way described in Bal-cells&Peletier(1994),Peletier&Balcells(1997)and de Grijs(1997).Both corrections are small(resp.≤0.10and≤0.04for I−K),because infrared colours of nearby galaxies are generally not very sensitive to these corrections.In order to compare these colours with central aper-ture measurements of ellipticals,a correction for internal colour gradients should be made.The colours of the sam-ple of Peletier&Balcells(1997)were taken on the minor axis in the bulge at0.5eﬀective bulge radii.Peletier&Bal-cells(1996)and also Terndrup et al.(1994)show that for their sample the colour diﬀerence in I−K between the bulge and the inner disk(2exponential scale lengths)is very small(0.07±0.15mag(Peletier&Balcells1996)),in-dicating small stellar population diﬀerences.For the edge-on galaxies of de Grijs et al.(1997)the colours were taken at an average height above the plane of2scale heights.Very little is known about vertical colour gradients in spirals.In our Galaxy Trefzger et al.(1995)ﬁnd a vertical metallicity gradient of∆[Fe/H]/∆h z=0.05,using h z=247pc(Kent et al.1991).Using simple stellar population models(e.g. Vazdekis et al.1996),this corresponds to∆(I−K)/∆h z =0.025.For external galaxies at present no data is avail-able for vertical stellar population gradients.Radial colour gradients are small as well-for the sample of Peletier& Balcells(1997)weﬁnd an average R−K gradient of0.12 mag per scale length,and a scatter of0.11mag.This corre-sponds to an I−K gradient of about0.10mag/scale length. Some simple modelling shows that we need to correct our colours by≤0.075mag to make them compatible with those of Bower et al.(1992a).Since we do not know how vertical colour gradients vary as a function of morphological type, we would have to apply the same correction for each galaxy, and therefore the slope of the CM relation in spirals would not change.For this reasonﬁnally the correction for internal stellar population gradients was not applied.3THE SCATTER IN THE CM RELATION, AND COMPARISON WITH EARLY-TYPEGALAXIESIn Fig.1we show the CM relation for our2samples.The drawn line is the least squaresﬁt.We have also plotted the elliptical and S0galaxies from Bower et al.(1992a),for which they found that the scatter was comparable with the obser-vational uncertainties.Since these authors do not have I−K colours,we have converted their V−K to I−K using the models of Vazdekis et al.(1996).To do so we have made a linearﬁt of I−K as a function of V−K for all single-age,single-metallicity models presented.Similar results are obtained if other stellar population models(e.g.,Worthey 1994)are used.Note that a sequence in morphological types can be seen in Fig.1,with the latest type spirals(ﬁlled circles)the faintest and the bluest.In Table1the best-ﬁtting least squaresﬁts are given. For our spiral galaxies we have applied a bivariateﬁtting routine,taking into account errors in both directions(see Peletier&Willner1993).The uncertainties in our photom-c 0000RAS,MNRAS000,000–000The scatter in the near-infrared colour-magnitude relation in spiral galaxies3Figure 1.The I −K vs M K relation as discussed in this paper.Plotted are the data from Peletier &Balcells (1997)with type earlier than 1.5(ﬁlled squares).The galaxies of de Grijs et al.(1997)are plotted as ﬁlled triangles for types earlier than 5.5and as crosses for later types.The elliptical/S0galaxies from Bower et al.(1992a)are shown as open circles.Table 1.CM relations (I −K =a +b (M K +25))and scat-ter for ellipticals and spirals.(1)gives the measured scatter in magnitudes,and (2)the scatter after correcting for observational uncertainties.intercept slope (1)(2)etry are considerably larger than for the ellipticals,because of the uncertainties in calibrating the infrared photometry with IRAC2b and the low light levels at large distances from the galactic planes4DISCUSSIONAlthough the scatter that we ﬁnd in the CM relation for spiral galaxies is much smaller than the values that have ap-peared in the literature until now,the results are consistent.Tully et al.(1982)already showed that spiral galaxies obey a CM relation,similar to the one found for ellipticals.They found that S0galaxies occupy a completely diﬀerent region on the (B T −H −0.5)CM plane than the later type spirals.This result,however,is peculiar for two reasons.First,some of the S0galaxies are redder than the bright elliptical galax-ies presented in the same paper.Balcells &Peletier (1994)showed that bulges of early-type spirals and S0s in U −R ,B −R and R −I are always bluer or have the same colour as ellipticals of the same luminosity,and therefore these S0galaxies must be aﬀected considerably by extinction (at least A B =0.5mag),which seems to be in conﬂict with the very small scatter among the S0galaxies.Secondly,there is a considerable gap of about 0.5mag between the S0galaxies and the early-type spirals,also unlikely from other work.Similarly,the relation of Wyse (1982)for spirals is probably severely aﬀected by extinction,since her brightest galaxies have B −H values around 5,much redder than the reddest nearby ellipticals without dust.In any case,much better in-terpretable and understandable is the work of Mobasher et al.(1986).They present CM relations in infrared (J −K )and optical-infrared (B −K )colours for spirals and ellipti-cals.In both bands the relation for spirals is steeper than that for ellipticals,although in the IR this diﬀerence is not very signiﬁcant.Their result,although with large scatter,is largely consistent with our results presented here.From Table 1one can see that the scatter in the CM relation for spirals is very small (0.07mag after taking into account the observational errors,or possibly less,since this number is smaller than the average observational error).Also,this relation falls about 0.1–0.4mag below the CM re-lation for ellipticals,and has a steeper slope.Here we present two possible explanations for the diﬀerence between the two CM relations.(i)While elliptical galaxies show very little or no star formation,and have at most 10%of their stars formed in the last 5Gyr (Bower et al.1992b),all spiral galaxies are currently forming stars,whereby the rate of star formation is determined only by the infrared luminosity (or mass)of the galaxy.(ii)Spiral galaxies in general have considerable vertical colour gradients.In this case we deduce that the vertical gra-dient per scale height needs to be about 0.15dex in [Fe/H],if the dust-free CM relation is the same for ellipticals and ter type spirals need to have larger gradients than earlier types.In the ﬁrst case,the most straightforward explanation for the origin of both relations would be that they are driven by metallicity,but that the bluer I −K colours are due to a small amount of recent star formation,that does not change the metallicity signiﬁcantly,but aﬀects the colours,and somewhat brightens the galaxy.In Fig.2we show a grid of Single Stellar Population (SSP)models of Bruzual &Charlot (1998)for metallicities between 0.004and solar and ages between 0.1and 20Gyr.We have assumed that the ellipticals are 20Gyr old and lie on the CM relation,and have followed this relation in time.Although the as-sumption that stars in a spiral galaxy are coeval is almost certainly too simple,the ﬁgure shows that on average the stars in late-type spirals are much younger than in ellipticals.The models here show ages ≤1Gyr,although one can ﬁnd other numbers if e.g.models with exponentially decreasing star formation are used.The most interesting aspect of this work is that we ﬁnd that the scatter in the CM relation for spirals is very small,and that there is a gap between spirals and ellipticals.This means that the current star formation in a spiral galaxy is determined by its size,morphological type or luminosity,and probably not by its environment or interactions,without much scatter.The CM relation can bec0000RAS,MNRAS 000,000–0004R.F.Peletier and R.deGrijsFigure 2.The I −K vs M K relation,now together with SSP models of Bruzual &Charlot (1998),for which metallicities and ages are given.better understood by comparing it to the large compilation of spiral galaxies by McGaugh &de Blok (1997),who show that fainter galaxies generally have lower surface brightness,bluer optical colours,and increasing gas mass fractions,as a consequence of which it is likely that their average ages are also younger.Now,let us assume that spiral galaxies have the same I −K CM relation as ellipticals,but that the diﬀerence we observe here is solely caused by vertical stellar population gradients in spiral galaxies.There are two reasons why we argue that this option is not very likely.First,the vertical metallicity gradient in our Galaxy,the only place where such a quantity has been measured at present,would correspond to only 0.025mag per scale height in I −K ,which is by far not enough to explain the diﬀerence between spirals and ellipticals.Secondly,the gradients would have to be larger for smaller galaxies,whereas the scatter between the spirals of the same luminosity would still need to remain very small.De Jong (1996)shows that radial colour gradients for the later types are not larger than those for larger,early-type spirals.For this reason it is not expected that the behavior in the vertical direction would be completely diﬀerent.5CONCLUSIONSThe main results of this paper are as follows:•We have determined a dust-free CM relation for spiral galaxies,by using I −K colours in edge-on galaxies above the plane.We ﬁnd that the scatter in this relation is small and approximately as large as can be explained by observational uncertainties.The slope of the IR CM-relation is steeper for spirals than for elliptical galaxies.•We present two possible explanations.First,the dif-ference could be caused by vertical colour gradients in spi-ral galaxies.In that case these gradients should be similar from galaxy to galaxy,have an average size of about 0.15dex in [Fe/H]per scale height,and increase for later galaxy types.The other,much more likely,possibility is that spirals and ellipticals have diﬀerent CM relations.The diﬀerence is caused by current star formation,which has to be present in all spirals,as opposed to ellipticals.The amount of current star formation would depend only on the galaxy’s infrared luminosity,and not on its environment.ACKNOWLEDGEMENTSR.de Grijs thanks the Dept.of Physics of the University of Durham for their hospitality during two visits.This paper is partly based on observations collected at the Isaac Newton Telescope,La Palma,the United Kingdom InfraRed Tele-scope,Mauna Kea,and the European Southern Observatory,La Silla.REFERENCESArimoto,N.,Yoshii,Y.,1987,A&A 173,23Balcells,M.,Peletier,R.F.,1994,AJ 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Spitzer-IRS mapping of the inner kpc of NGC 253 Spatial distribution of the Ne III, PAH 11.

and a gradient in the [Ne iii]/[Ne ii] line ratio.
Daniel Devost1, Bernhard R. Brandl1,2, L. Armus3, D. J. Barry1, G.C. Sloan1, Vassilis Charmandaris1, Henrik Spoon1, Jeronimo Bernard-Salas1, and James R. Houck1
arXivБайду номын сангаасastro-ph/0406172v1 7 Jun 2004
Spitzer-IRS mapping of the inner kpc of NGC 253: Spatial distribution of the [Ne iii], PAH 11.3 µm and H2 (0-0) S(1) lines
Four luminous super star clusters were discovered in the central region by Watson et al. (1996) with the Hubble Space Telescope. They derived a bolometric luminosity for the brightest cluster of 1.3 × 109L⊙ which corresponds to 10% of the luminosity of within a region of a radius of 15′′ at the nucleus. Also, observations of Hα emission and earlier Einstein and ROSAT x-ray data (see Ptak et al. 1997, and references therein) revealed a starburst-driven wind, emanating from the nucleus along the minor axis of the galaxy. This wind was well delineated by Chandra (Strickland, Heckman, Weaver, & Dahlem 2000). Engelbracht et al. (1998) studied NGC 253 using near-infrared and mid-infrared lines and concluded that the properties of the Initial Mass Function (IMF) in the starbursting region is similar to a Miller-Scalo IMF that has a deﬁciency in low mass stars. Also, Boeker, Krabbe, & Storey (1998) and Keto et al. (1999) presented maps of [Ne ii] 12.81 µm and PAH emission.

火山作文英语

Volcanoes are one of the most fascinating and powerful natural phenomena on Earth. They are geological formations that result from the eruption of molten rock,known as magma,from beneath the Earths surface.Here is a detailed exploration of volcanoes in an English essay format:IntroductionA volcano is a vent in the Earths crust through which molten rock,gases,and ash are expelled.They are often associated with the destructive power they can unleash,but they also play a crucial role in the formation of new land and the recycling of the Earths crust.Types of Volcanoes1.Cinder Cone Volcanoes:These are the simplest type of volcano,formed from the accumulation of small,solidified lava fragments that have been ejected from the vent.posite Volcanoes:Also known as stratovolcanoes,they are made up of alternating layers of lava and ash,often with a steep profile and a summit crater.3.Shield Volcanoes:These have a broad,gently sloping profile,resembling a warriors shield.They are formed from the eruption of lowviscosity lava that flows easily.4.Caldera Volcanoes:Formed by the collapse of a volcanic structure following a major eruption,they often have a large,bowlshaped depression.Formation of VolcanoesVolcanoes form at plate boundaries where the Earths crust is either spreading apart or colliding.At divergent boundaries,magma rises to the surface,creating new crust and often shield volcanoes.At convergent boundaries,one plate is forced under another, melting and creating magma that can rise to form composite volcanoes.Eruption ProcessA volcanic eruption begins with seismic activity and ground deformation as magma rises towards the surface.This is often followed by the ejection of gases,ash,and lava.The type of eruption can vary from quiet effusive eruptions to explosive ones,which can cause widespread devastation.Effects of Volcanic Eruptions1.Environmental Impact:Volcanic ash can disrupt air travel,affect climate by reflecting sunlight,and enrich soil with minerals.2.Geological Impact:Lava flows can create new land and reshape existing landscapes.3.Human Impact:While they can cause loss of life and property,they also provide fertile soil for agriculture and geothermal energy resources.Volcano Monitoring and SafetyScientists monitor volcanoes using seismographs,GPS,and satellite imagery to predict eruptions and mitigate risks.Evacuation plans and early warning systems are crucial for communities living near active volcanoes.ConclusionVolcanoes are aweinspiring features of our planet,offering both beauty and danger. Understanding their formation,types,and the processes involved in eruptions is essential for appreciating their role in the Earths geological history and for ensuring the safety of those living in their shadow.。

A Submillimeter and Radio Survey of Gamma-Ray Burst Host Galaxies A Glimpse into the Future

a rXiv:as tr o-ph/21645v21Nov22A Submillimeter and Radio Survey of Gamma-Ray Burst Host Galaxies:A Glimpse into the Future of Star Formation Studies E.Berger 1,L.L.Cowie 2,S.R.Kulkarni 1,D.A.Frail 3,H.Aussel 2,&A.J.Barger 2,4,5ABSTRACT We present the ﬁrst comprehensive search for submillimeter and radio emis-sion from the host galaxies of twenty well-localized γ-ray bursts (GRBs).With the exception of a single source,all observations were undertaken months to years after the GRB explosions to ensure negligible contamination from the afterglows.We detect the host galaxy of GRB 000418in both the sub-mm and radio,and the host galaxy of GRB 000210only in the sub-mm.These observations,in con-junction with the previous detections of the host galaxies of GRB 980703and GRB 010222,indicate that about 20%of GRB host galaxies are ultra-luminous (L >1012L ⊙)and have star formation rates of about 500M ⊙yr −1.As an en-semble,the non-detected hosts have a star formation rate of about 100M ⊙yr −1(5σ)based on their radio emission.The detected and ensemble star formation rates exceed the optically-determined values by an order of magnitude,indicat-ing signiﬁcant dust obscuration.In the same vein,the ratio of bolometric dust luminosity to UV luminosity for the hosts detected in the sub-mm and radio ranges from ∼20−800,and follows the known trend of increasing obscuration with increasing bolometric luminosity.We also show that,both as a sample and individually,the GRB host galaxies have bluer R −K colors as compared with galaxies selected in the sub-mm in the same redshift range.This possibly in-dicates that the stellar populations in the GRB hosts are on average younger,supporting the massive stellar progenitor scenario for GRBs,but it is also possible that GRB hosts are on average less dusty.Beyond the speciﬁc results presentedin this paper,the sub-mm and radio observations serve as an observational proof-of-concept in anticipation of the upcoming launch of the SWIFT GRB missionand SIRTF.These new facilities will possibly bring GRB host galaxies into theforefront of star formation studies.Subject headings:cosmology:observations—galaxies:starburst—gamma-rays:bursts—stars:formation1.IntroductionOne of the major thrusts in modern cosmology is an accurate census of star formation and star-forming galaxies in the Universe.This endeavor forms the backbone for a slew of methods(observational,analytical,and numerical)to study the process of galaxy formation and evolution over cosmic time.To date,star-forming galaxies have been selected and studied mainly in two observational windows:the rest-frame ultraviolet(UV),and rest-frame radio and far-infrared(FIR).For galaxies at high redshift these bands are shifted into the optical and radio/sub-mm,allowing observations from the ground.Still,the problem of translating the observed emission to star formation rate(SFR)involves large uncertainty.This is partly because each band traces only a minor portion of the total energy output of stars.Moreover, the optical/UV band is signiﬁcantly aﬀected by dust obscuration,thus requiring order of magnitude corrections,while the sub-mm and radio bands lack sensitivity,and therefore uncover only the most prodigiously star-forming galaxies.The main result that has emerged from star formation surveys over the past few years is exempliﬁed in the so-called Madau ly,the SFR volume density,ρSFR(z), rises steeply to z∼1,and seemingly peaks at z∼1−2.There is still some debate about the how steep the rise is,with values ranging from(1+z)1.5(Wilson et al.2002)to(1+z)4 (e.g.Madau et al.1996).The evolution beyond z∼2is even less clear since optical/UV observations indicate a decline(Madau et al.1996),while recent sub-mm observations argue for aﬂatρSFR(z)to higher redshift,z∼4(Barger,Cowie&Richards2000).Consistency with this trend can be obtained by invoking large dust corrections in the optical/UV(Steidel et al.1999).For general reviews of star formation surveys we refer the reader to Kennicut (1998),Adelberger&Steidel(2000),and Blain et al.(2002).Despite the signiﬁcant progress in thisﬁeld,our current understanding of star formation and its redshift evolution is still limited by the biases and shortcomings of current optical/UV, sub-mm,and radio selection techniques.In particular,despite the fact that the optical/UV band is sensitive to galaxies with modest star formation rates(down to a fraction of a M⊙yr−1)at high redshift,these surveys may miss the most dusty,and vigorously star-forming galaxies.Moreover,it is not clear if the simple,locally-calibrated prescriptions for correcting the observed un-obscured SFR for dust extinction(e.g.Meurer,Heckman&Calzetti1999), hold at high redshift;even if they do,these prescriptions involve an order of magnitude correction.Finally,the optical/UV surveys are magnitude limited,and therefore miss the faintest sources.Sub-mm surveys have uncovered a population of highly dust-extincted galaxies,which are usually optically faint,and have star formation rates of several hundred M⊙yr−1(e.g.Smail, Ivison&Blain1997).However,unlike optical/UV surveys,sub-mm surveys are severely sen-sitivity limited,and only detect galaxies with L bol 1012L⊙.More importantly,current sub-mm bolometer arrays(such as SCUBA)have large beams on the sky(∼15arcsec) making it diﬃcult to unambiguously identify optical counterparts(which are usually faint to begin with),and hence measure the redshifts(Smail et al.2002);in fact,of the∼200 sub-mm galaxies identiﬁed to date,only a handful have a measured redshift.Finally,trans-lating the observed sub-mm emission to a SFR requires signiﬁcant assumptions about the temperature of the dust,and the dust emission spectrum(e.g.Blain et al.2002).Surveys at decimeter radio wavelengths also suﬀer from low sensitivity,but the high as-trometric accuracy aﬀorded by synthesis arrays such as the VLA allows a sub-arcsec localiza-tion of the radio-selected galaxies.As a result,it is easier to identify the optical counterparts of these sources.Recently,this approach has been used to pre-select sources for targeted sub-mm observations resulting in an increase in the sub-mm detection rate(Barger,Cowie &Richards2000;Chapman et al.2002)and redshift determination(Chapman et al.2003). However,this method is biased towardﬁnding luminous(high SFR)sources since it requires an initial radio detection.An additional problem with radio,even more than with sub-mm, selection is contamination by active galactic nuclei(AGN).An examination of the X-ray properties of radio and sub-mm selected galaxies reveals that of the order of20%can have a signiﬁcant AGN component(Barger et al.2001).The most signiﬁcant problem with current star formation studies,however,is that the link between the optical and sub-mm/radio samples is still not well understood.The Hubble Deep Field provides a clear illustration:the brightest sub-mm source does not appear to have an optical counterpart(Hughes et al.1998),and only recently a detection has been claimed in the near-IR(K≈23.5mag;Dunlop et al.2002).Along the same line,sub-mm observations of the optically-selected Lyman break galaxies have resulted in very few detections(Chapman et al.2000;Peacock et al.2000;Chapman et al.2002),and even the brightest Lyman break galaxies appear to be faint in the sub-mm band(Baker et al. 2001).In addition,there is considerable diversity in the properties of optical counterpartsto sub-mm sources,ranging from galaxies which are faint in both the optical and near-IR (NIR)to those which are bright in both bands(Ivison et al.2000;Smail et al.2002).As a result of the unclear overlap,and the sensitivity and dust problems in the sub-mm and optical surveys,the fractions of global star formation in the optical and sub-mm/radio selected galaxies is not well constrained.It is therefore not clear if the majority of star formation takes place in ultra-luminous galaxies with very high star formation rates,or in the more abundant lower luminosity galaxies with star formation rates of a few M⊙yr−1. Given the diﬃculty with redshift identiﬁcation of sub-mm galaxies,the redshift distribution of dusty star forming galaxies remains highly uncertain.One way to alleviate some of these problems is to study a sample of galaxies that is immune to the selection biases of current optical/UV and sub-mm/radio surveys,and which may draw a more representative sample of the underlying distribution of star-forming galaxies.The host galaxies ofγ-ray bursts(GRBs)may provide one such sample.The main advantages of the sample of GRB host galaxies are:(i)The galaxies are selected with no regard to their emission properties in any wavelength regime,(ii)the dust-penetrating power of theγ-ray emission results in a sample that is completely unbiased with respect to the global dust properties of the hosts,(iii)GRBs can be observed to very high redshifts with existing missions(z 10;Lamb&Reichart2000),and as a result volume corrections for the star formation rates inferred from their hosts are trivial,(iv)the redshift of the galaxy can be determined via absorption spectroscopy of the optical afterglow,or X-ray spectroscopy allowing a redshift measurement of arbitrarily faint galaxies(the current record-holder is the host of GRB990510with R=28.5mag and z=1.619;Vreeswijk et al.2001),and(v)since there is excellent circumstantial evidence linking GRBs to massive stars(e.g.Bloom,Kulkarni&Djorgovski2002,the sample of GRB hosts is expected to trace global star formation(Blain&Natarajan2000).Of course,the sample of GRB hosts is not immune from its own problems and potential biases.The main problem is the relatively small size of the sample in comparison to both the optical and sub-mm samples6(although the number of GRB hosts with a known redshift exceeds the number of sub-mm galaxies with a measured redshift).As a result,at the present it is not possible to assess the SFR density that is implied by GRB hosts,or its redshift evolution.A bias towards sub-solar metallicity for GRB progenitors(and hence their environments)has been discussed(MacFadyen&Woosley1999;MacFadyen,Woosley &Heger2001),but it appears that very massive stars(e.g.M 35M⊙)should produceblack holes even at solar metallicity.The impact of metallicity on additional aspects of GRB formation(e.g.angular momentum,loss of hydrogen envelope)is not clear at present. Finally,given the observed dispersion in metallicity within galaxies(e.g.Alard2001;Overzier et al.2001),it is likely that even if GRBs require low metallicity progenitors,this does not imply that the galaxy as a whole has a lower than average metallicity.To date,GRB host galaxies have mainly been studied in the optical and NIR bands. With the exception of one source(GRB020124;Berger et al.2002),every GRB localized to a sub-arcsecond position has been associated with a star-forming galaxy(Bloom,Kulkarni &Djorgovski2002).These galaxies range from R≈22−29mag,have a median redshift, z ∼1,and are generally typical of star-forming galaxies at similar redshifts in terms of morphology and luminosity(Djorgovski et al.2001),with star formation rates from optical spectroscopy of∼1−10M⊙yr−1.At the same time,there are hints for higher than average ratios of[Ne III]3869to[O II]3727,possibly indicating the presence of massive stars (Djorgovski et al.2001).Only two host galaxies have been detected so far in the radio (GRB980703;Berger,Kulkarni&Frail2001)and sub-mm(GRB010222;Frail et al.2002).Here we present sub-mm and radio observations of a sample of20GRB host galaxies, ranging in redshift from about0.4to4.5(§2);one of the20sources is detected with high signiﬁcance in both the sub-mm and radio bands,and an additional source is detected in the sub-mm(§3).We compare the detected sub-mm and radio host galaxies to local and high-z ultra-luminous galaxies in§4,and derive the SFRs in§5.We then compare the inferred SFRs of the detected host galaxies,and the ensemble of undetected hosts,to optical estimates in §6.Finally,we compare the optical properties of the GRB host galaxies to those of sub-mm and radio selected star-forming galaxies(§7).2.Observations2.1.Target SelectionAt the time we conducted our survey,the sample of GRB host galaxies numbered25, twenty of which had measured redshifts.These host galaxies were localized primarily based on optical afterglows,but also using the radio and X-ray afterglow emission.Of the25host galaxies we observed eight in both the radio and sub-mm,seven in the radio,andﬁve in the sub-mm.The galaxies were drawn from the list of25hosts at random,constrained primarily by the availability of observing time.Thus,the sample presented here does not suﬀer from any obvious selection biases,with the exception of detectable afterglow emission in at least one band.Sub-mm observations of GRB afterglows,and a small number of host galaxies have been undertaken in the past.Starting in1997,Smith et al.(1999)and Smith et al.(2001)have searched for sub-mm emission from the afterglow of thirteen GRBs.While they did not detect any afterglow emission,these authors used their observations to place constraints on emission from eight host galaxies,with typical1σrms values of1.2mJy.Since these were target-of-opportunity observations,they were not always undertaken in favorable observing conditions.More recently,Barnard et al.(2002)reported targeted sub-mm observations of the host galaxies of four optically-dark GRBs(i.e.GRBs lacking an optical afterglow).They focused on these particular sources since one explanation for the lack of optical emission is obscuration by dust,which presumably points to a dusty host.None of the hosts were detected,with the possible exception of GRB000210(see§3.4),leading the authors to conclude that the hosts of dark bursts are not necessarily heavily dust obscured.Thus,the observations presented here provide the most comprehensive and bias-free search for sub-mm emission from GRB host galaxies,and theﬁrst comprehensive search for radio emission.2.2.Submillimeter ObservationsObservations in the sub-mm band were carried out using the Sub-millimeter Common User Bolometer Array(SCUBA;Holland et al.1999)on the James Clerk Maxwell Tele-scope(JCMT7).We observed the positions of thirteen well-localized GRB afterglows with the long(850µm)and short(450µm)arrays.The observations,summarized in Table1, were conducted in photometry mode with the standard nine-jiggle pattern using the central bolometer in each of the two arrays to observe the source.In the case of GRB000301C we used an oﬀ-center bolometer in each array due to high noise levels in the central bolometer.To account for variations in the sky brightness,we used a standard chopping of the secondary mirror between the on-source position and a position60arcsec away in azimuth, at a frequency of7.8125Hz.In addition,we used a two-position beam switch(nodding), in which the beam is moved oﬀ-source in each exposure to measure the sky.Measurements of the sky opacity(sky-dips)were taken approximately every two hours,and the focus andarray noise were checked at least twice during each shift.The pointing was checked approximately once per hour using several sources throughout each shift,and was generally found to vary by 3arcsec(i.e.less than one quarter of the beam size).All observations were performed in band2and3weather withτ225GHz≈0.05−0.12.The data were initially reduced with the SCUBA Data Reduction Facility(SURF)fol-lowing the standard reduction procedure.The oﬀ-position pointings were subtracted from the on-position pointings to account for chopping and nodding of the telescope.Noisy bolometers were removed to facilitate a more accurate sky subtraction(see below),and the data were thenﬂat-ﬁelded to account for the small diﬀerences in bolometer response.Ex-tinction correction was performed using a linear interpolation between skydips taken before and after each set of on-source scans.In addition to the sky subtraction oﬀered by the nodding and chopping,short-term sky contributions were subtracted by using all low-noise oﬀ-source bolometers(sky bolometers). This procedure takes advantage of the fact that the sky contribution is correlated across the array.As a result,theﬂux in the sky bolometers can be used to assess the sky contribution to theﬂux in the on-source bolometer.This procedure is especially crucial when observing weak sources,since the measuredﬂux may be dominated by the sky.We implemented the sky subtraction using SURF and our own routine using MATLAB.We found that in general the SURF sky subtraction under-estimated the sky contribution,and as a result over-estimated the sourceﬂuxes.We therefore used the results of our own analysis routine.For this purpose we calculated the median value of the two(three)outer rings of bolometers in the850µm (450µm)array,after removing noisy bolometers(deﬁned as those whose standard deviation over a whole scan deviated by more than5σfrom the median standard deviation of all sky bolometers).Following the sky subtraction,we calculated the mean and standard deviation of the mean(SDOM)for each source in a given observing shift.Noisy data were eliminated in two ways.First,the data were binned into25equal time bins,and the SDOM was calculated step-wise,i.e.at each step the data from an additional bin were added and the mean and SDOM were re-calculated.In an ideal situation where the data quality remains approximately constant,the SDOM should progressively decrease as more data are accumulated.However, if the quality of the data worsens(due to deteriorating weather conditions for example) the SDOM will increase.We therefore removed time bins in which the SDOM increased. Following this procedure,we recursively eliminated individual noisy data points using a sigma cutoﬀlevel based on the number of data points(Chauvenet’s criterion;Taylor1982)until the mean converged on a constant value.Typically,two or three iterations were required,with only a few data points rejected each time.Typically,only a few percent of the data were rejected by the two procedures.Finally,ﬂux conversion factors(FCFs)were applied to the resulting voltage measure-ments to convert the signal to ing photometry observations of Mars and Uranus,and/or secondary calibrators(OH231.8+4.2,IRC+10216,and CRL618),we found the FCFto vary between180−205Jy/V at850µm,consistent with the typical value of197±13.At450µm,the FCFs varied between250−450Jy/V.2.3.Radio ObservationsVery Large Array(VLA8):We observed12GRB afterglow positions with the VLA from April2001to February2002.All sources were observed at8.46GHz in the standard continuum mode with2×50MHz bands.In addition,GRB000418was observed at1.43and4.86GHz,and GRB0010222was observed at4.86GHz.In Table2we provide a summaryof the observations.In principle,since the median spectrum of faint radio sources between1.4and8.5GHzis Fν∝ν−0.6(Fomalont et al.2002),the ideal VLA frequency for our observations(takinginto account the sensitivity at each frequency)is1.43GHz.However,we chose to observe primarily at8.46GHz for the following reason.The majority of our observations were takenin the BnC,C,CnD,and D conﬁgurations,in which the typical synthesized beam size is∼10−40arcsec at1.43GHz,compared to∼2−8arcsec at8.46GHz.The large synthesized beam at1.43GHz,combined with the largerﬁeld of view and higher intrinsic brightnessof radio sources at this frequency,would result in a signiﬁcant decrease in sensitivity dueto source confusion.Thus,we were forced to observe at higher frequencies,in which the reduced confusion noise more than compensates for the typical steep spectrum.We chose8.46GHz rather than4.86GHz since the combination of20%higher sensitivity and60% lower confusion noise,provide a more signiﬁcant impact than the typical30%decrease in intrinsic brightness.The1.43GHz observations of GRB000418were undertaken in the Aconﬁguration,where confusion does not play a limiting role.Forﬂux calibration we used the extragalactic sources3C48(J0137+331),3C147(J0542+498), and3C286(J1331+305),while the phases were monitored using calibrator sources within∼5◦of the survey sources.We used the Astronomical Image Processing System(AIPS)for data reduction and analysis.For each source we co-added all the observations prior to producing an image,to increase theﬁnal signal-to-noise.Australia Telescope Compact Array(ATCA9):We observed the positions of four GRB afterglows during April2002,in the6A conﬁguration ing the6-km baseline resulted in a signiﬁcant decrease in confusion noise,thus allowing observations at the most advantageous frequencies.The observations are summarized in Table2.We used J1934−638forﬂux calibration,while the phase was monitored using calibrator sources within∼5◦of the survey sources.The data were reduced and analyzed using the Multichannel Image Reconstruction,Image Analysis and Display(MIRIAD)package,and AIPS.3.ResultsTheﬂux measurements at the position of each GRB are given in Tables1and2,and are plotted in Figure2.Of the20sources,only GRB000418was detected in both the radio and sub-mm with S/N>3(§3.1).One additional source,GRB000210,is detected with S/N>3 when combining our observations with those of Barnard et al.(2002).Two hosts have radio ﬂuxes with3<S/N<4(GRB000301C and GRB000926),but as we show below this is due in part to emission from the afterglow.The typical2σthresholds are about2mJy,20µJy,and70µJy in the SCUBA,VLA,and ATCA observations,respectively.In Figure2we plot all sources with S/N>3as detections, and the rest as2σupper limits.In addition,for the sources observed with the ATCA we plot both the1.4GHz upper limits,and the inferred upper limits at8.46GHz assuming a typical radio spectrum,Fν∝ν−0.6(Fomalont et al.2002).One obvious source for the observed radio and sub-mmﬂuxes(other than the putative host galaxies)is emission from the afterglows.To assess the possibility that our observa-tions are contaminated byﬂux from the afterglows we note that the observations have been undertaken at least a year after the GRB explosion10.On this timescale,the sub-mm emis-sion from the afterglow is expected to be much lower than the detection threshold of ourobservations.In fact,the brightest sub-mm afterglows to date have only reached aﬂux of a few mJy(at350GHz),and typically exhibited a fading rate of∼t−1after about one day following the burst(Smith et al.1999;Berger et al.2000;Smith et al.2001;Frail et al. 2002;Yost et al.2002).Thus,on the timescale of our observations,the expected sub-mm ﬂux from the afterglows is only∼10µJy,well below the detection threshold.The radio emission from GRB afterglows is more long-lived,and hence posses a more serious problem.However,on the typical timescale of the radio observations the8.46GHz ﬂux is expected to be at most a fewµJy(e.g.Berger et al.2000).In the following sections we discuss the individual detections in the radio and sub-mm, and also provide an estimate for the radio emission from each afterglow.3.1.GRB000418A source at the position of GRB000418is detected at four of theﬁve observing fre-quencies with S/N>3.The SCUBA source,which we designate SMM12252+2006,has a ﬂux density of Fν(350GHz)≈3.2±0.9mJy,and Fν(670GHz)≈41±19mJy.These values(Fν∝νβ),consistent with a thermal dust spectrum as imply a spectral index,β≈3.9+1.1−1.3expected if the emission is due to obscured star formation.The radio source(VLA122519.26+200611.1),is located atα(J2000)=12h25m19.255s,δ(J2000)=20◦06′11.10′′,with an uncertainty of0.1arcsec in both coordinates.This position is oﬀset from the position of the radio afterglow of GRB000418(Berger et al.2001)by ∆α=−0.40±0.14arcsec and∆δ=−0.04±0.17arcsec(Figure1).In comparison,the oﬀset measured from Keck and Hubble Space Telescope images is smaller,∆α=−0.019±0.066 arcsec and∆δ=0.012±0.058arcsec.VLA122519.26+200611.1has an observed spectral slopeβ=−0.17±0.25,ﬂatter than the typical value for faint radio galaxies,β≈−0.6(Fomalont et al.2002),and similar to the value measured for the host of GRB980703(β≈−0.32;Berger,Kulkarni&Frail2001). The source appears to be slightly extended at1.43and8.46GHz,with a size of about1 arcsec(8.8kpc at z=1.119).The expected afterglowﬂuxes at4.86and8.46GHz at the time of our observations are about5and10µJy,respectively(Berger et al.2001).At1.43GHz the afterglow contribution is expected to be about10µJy based on the4.86GHzﬂux and the afterglow spectrum Fν∝ν−0.65.Thus,despite the contribution from the afterglow,the radio detections of the host galaxy are still signiﬁcant at better than3σlevel.Correcting for the afterglowcontribution weﬁnd an actual spectral slopeβ=−0.29±0.33,consistent with the median β≈−0.6for8.46GHz radio sources with a similarﬂux(Fomalont et al.2002).As with all SCUBA detections,source confusion arising from the large beam(D FWHM≈14arcsec at350GHz and≈6arcsec at670GHz)raises the possibility that SMM12252+2006 is not associated with the host galaxy of GRB000418.Fortunately,the detection of the radio source,which is located0.4±0.1arcsec away from the position of the radio afterglow of GRB000418,indicates that SMM12252+2006and VLA122519.26+200611.1are in fact the same source—the host galaxy of GRB000418.Besides the positional coincidence of the VLA and SCUBA sources,we gain further.This conﬁdence of the association based on the spectral index between the two bands,β3501.4 spectral index is redshift dependent as a result of the diﬀerent spectral slopes in the two regimes(Carilli&Yun2000;Barger,Cowie&Richards2000).Weﬁndβ350≈0.73±0.10,1.4=0.59±0.16(for the redshift in good agreement with the Carilli&Yun(2000)value ofβ3501.4of GRB000418,z=1.119).We also detect another source,slightly extended(θ≈1arcsec),approximately1.4arcsec East and2.7arcsec South of the host of GRB000418(designated VLA122519.36+200608.4), with Fν(1.43GHz)=48±15µJy and Fν(8.46GHz)=37±12µJy(see Figure1).This source appears to be linked by a bridge of radio emission(with S/N≈1.5at both frequencies)to the host of GRB000418.The physical separation between the two sources,assuming both are at the same redshift,z=1.119,is25kpc.There is no obvious optical counterpart to this source in Hubble Space Telescope images down to about R∼27.5mag.Based purely on radio source counts at8.46GHz(Fomalont et al.2002),the expected number of sources with Fν(8.46GHz)≈37µJy in a3arcsec radius circle is only about 2.7×10−4.Thus,the coincidence of two such faint sources within3arcsec is highly suggestive of an interacting system,rather than chance superposition.Interacting radio galaxies with separations of about20kpc,and joined by a bridge of radio continuum emission have been observed locally(Condon et al.1993;Condon,Helou& Jarrett2002).In addition,optical surveys(e.g.Patton et al.2002)show that a few percent of galaxies with an absolute B-band magnitude similar to that of the host of GRB000418, have companions within about30kpc.The fraction of interacting systems is possibly much higher,∼50%,in ultra-luminous systems(such as the host of GRB000418),both locally (Sanders&Mirabel1996)and at high redshift(e.g.Ivison et al.2000).We note that with a separation of only3arcsec,the host of GRB000418and the companion galaxy fall within the SCUBA beam.Thus,it is possible that SMM12252+2006to about is in fact a superposition of both radio sources.This will change the value ofβ3501.4。

OSRAM OSTAR Observation应用注意事项说明书

April 8, 2011page 1 of 10OSRAM OSTAR Observation Application NoteSummaryThis application note provides an overview of the general handling and functionality of the OSRAM OSTAR Observation. The im-portant optical and electrical characteristics are described and the thermal requirements for stable operation of the IR LED light source are addressed.In addition, the procedure for dimensioning an appropriate heat sink is illustrated by means of an example.Applications of the IR light source OSRAM OSTAR ObservationThere are various possibilities where our customers are using the OSRAM OSTAR Observation as IR light source:- Infrared illumination for cameras - General monitoring systems - IR data transfer- Driver assistance systems.Due to its compact and flat design together with its high light density, the OSRAM OSTAR Observation can be easily inte-grated in various applications. This opens up new application areas that were off limits to conventional IR devices.Construction of the OSRAM OSTAR ObservationDuring design of the OSRAM OSTAR Ob-servation, special attention was given to the thermal optimization of the module.The module core is formed from ten highly efficient semiconductor chips mounted on ceramic. For optimal heat transfer, the ce-ramic is directly mounted to the aluminum of the insulated metal core circuit board (base plate). This results in optimal heat dissipa-tion and additionally provides a sufficiently large area for a good thermal connection to the system heat sink where the OSRAM OSTAR module has to be attached to.With this construction, the light source itself exhibits a very low thermal resistance (R thJB ) between junction and base plate of 2.8 K/W.The frame surrounding the chips is available in black and white colour to enable a choice depending on the desired application.The black frame minimizes scattered light, which is important in imaging systems, whereas the white frame optimizes the total optical output power.Figure 1: Two frame colours are available for the OSRAM OSTAR Observation.April 8, 2011page 2 of 10Equipped with an ESD protection diode, the OSRAM OSTAR Observation possesses ESD protection up to 2 kV according to JESD22-A114-B.A thermistor (NTC EPCOS 8502) mounted to the base plate serves as a sensor for de-termining the temperature of the metal core board. The NTC temperature provides a good approximation of the average tempera-ture of the underside of the aluminum base plate. From this the junction temperature can be estimated (using R thJB ) and thus con-trolled.As a light source, semiconductors of the latest highly efficient thin film technology based on AlGaAs are employed. This pro-vides a nearly pure surface emitter with Lambertian radiation characteristics.All semiconductor chips are wired in series to achieve a constant intensity for all emit-ting surfaces.Tips for handling the OSRAM OSTAR ObservationIn order to protect the semiconductor chips from environmental influences such as mois-ture, they are encapsulated using a clear silicone.In addition, the silicone encapsulant allows an operation at a junction temperature of 145°C.Since this encapsulant is very elastic and soft, mechanical damage to the silicone should be minimized or avoided if at all pos-sible during processing (see also the appli-cation note "Handling of Silicone Resin LEDs“).This also applies to the black silicone en-capsulant for the connection contacts. Ex-cessive force on the cover can lead to spon-taneous failure of the light source (damageto the contacts).Figure 2: Areas of the silicone encapsu-lant of the OSRAM OSTAR Observation (shown in red hatch marks), which must not be damaged.In Figure 2, the corresponding locations are shown in red hatch marks.To prevent damaging or puncturing the en-capsulant the use of all types of sharp ob-jects should be avoided.Furthermore, it should be assured that the light source is provided with adequate cool-ing (see design example below) during op-eration. Even at low currents, prolonged operation without cooling can lead to over-heating, damage or even failure of the mod-ule.Electrical connection of the OS-RAM OSTAR ObservationFor easy electrical connection, the OSRAM OSTAR Observation is equipped with a 4-pin socket:Pin Assignment: Pin 1: Anode Pin 2: Thermistor Pin 3: Thermistor Pin 4: CathodeAs a mating plug, the SMD plug from ERNI (SMD214025.4-pins) is recommended.April 8, 2011page 3 of 10Mounting the OSRAM OSTAR Ob-servationSeveral mounting methods can be used for attaching the IR light source.When selecting an appropriate mounting method, make sure that a good heat transfer is provided between the OSRAM OSTAR Observation and the heat sink and that this is also guaranteed during operation.An insufficient or incorrect mounting can lead to thermal or mechanical problems dur-ing assembly.Generally, screws should be used for mount-ing the OSRAM OSTAR Observation.When mounting the module with M2 screws, a torque of 0.2 - 0.3 Nm should be used. In order to achieve a good thermal connection, the contact pressure should typically be in the range of 0.35 MPa.In addition to mounting with screws, the OSRAM OSTAR Observation can also be attached by means of gluing or clamping. When mounting with glue, care should be taken that the glue is both adhesive and thermally stable, and possesses a good thermal conductivity.When mounting a component to a heat sink, it should generally be kept in mind that the two solid surfaces must be brought into physical contact.Technical surfaces are never really flat or smooth, however, but have a certain rough-ness due to microscopic edges and depres-sions. When two such surfaces are joined together, contact occurs only at the surface peaks. The depressions remain separated and form air-filled cavities (Figure 3).DescriptionMaterial Advantages DisadvantagesThermally conductive pasteTypically silicone based, with heat conductiveparticlesThermally conductive compoundsImproved thermallyconductive paste – rub-bery film after curingThinnest connection with minimal pressureHigh thermal conductiv-ity No delaminationMaterial discharge at the edgesDanger of contamina-tion during mass pro-ductionPaste can escape and "creep" over timeConnections require curing process Phase Change Materi-als (PCM)Material of polyester or acrylic with lower glass transition temperature, filled with thermally con-ductive particlesEasy handling and mountingNo delaminationNo curingContact pressure re-quiredHeat pretreatment re-quiredThermally conductive elastomersSilicone plastic washer pads- filled with thermally conductive particles - often strengthened with glass fibers or di-electric filmsThermally conductive tapeDouble sided tape filled with particles for uniform thermal and adhesive propertiesNo leakage of materialCuring not requiredProblem with delamina-tionModerate thermal con-ductivityContact pressure re-quiredTable 1: Thermal Interface MaterialsApril 8, 2011page 4 of 10Figure 3: Heat flow with and without heat conductive material.Since air is a poor conductor of heat, these cavities should be filled with a thermally conductive material in order to significantly reduce the thermal resistance and improve the heat flow between the two adjacent sur-faces.Without an appropriate, optimally effective interface, only a limited amount of heat ex-change occurs between the two surfaces, eventually leading to overheating of the light source.To improve the heat transfer capability and reduce the thermal contact resistance, sev-eral materials are suitable.Thermally conductive pastes and com-pounds possess the lowest transfer resis-tance, but require a certain amount of care in handling.Elastomers and foils/bands are easy to use. With pretreated surfaces and appropriate contact pressure, a good thermal transfer can be realized.Table 1 shows an overview of the most commonly used thermally conductive mate-rials along with their most important advan-tages and disadvantages.Optical characteristics of the OS-RAM OSTAR ObservationWhen characterizing IR LEDs, the intensity is usually specified with two parameters - the total radiant flux Φe (units of mW) and the radiant intensity I e (units of mW/sr).The total radiant flux Φe of an LED describes the total radiated light power independent of direction. For the OSRAM OSTAR Observa-tion, this is shown in Figure 4, in relation to forward current.In contrast, the radiant intensity expresses the radiated power within a fixed solid angle (e.g. 0.01 sr ≙ ±3.2°) in the primary direction of radiation (optical axis).Figure 4: Relative total radiant flux in re-lation to forward current I F .The radiation characteristics (in the far field ) show the distribution of intensity dependent on angle and are shown for the OSRAM OSTAR Observation in Figure 5. This repre-sents a good approximation of a Lambertian source with a radiation angle of ±60°.In general, the brightness can be influenced with the help of appropriate secondary op-tics. That is, with the use of focusing optics, the light output within a particular angle can be significantly increased.April 8, 2011page 5 of 10Figure 5: Radiation characteristics with-out optics.The user should refrain from attempting to mount the primary optics to the silicone en-capsulant. This can lead to damage to the chip and especially to the bonding wires, thereby voiding the warranty provided by OSRAM.In the near field (at different operating cur-rents), the OSRAM OSTAR Observation exhibits the radiance images shown in Fig-ure 6.Figure 6: Radiance images in the near field at very low power (above) and at higher power (below).An especially homogeneous radiance is achieved through the black frame of the module - a particular advantage when using imaging optics.Optical safety regulationsDepending on the mode of operation, the OSRAM OSTAR Observation emits highly concentrated, invisible infrared radiation, which can be dangerous for the human eye. Products which contain these components must be handled according to the guidelines specified in IEC Standard 60825-1 and IEC 62471 "Photobiological Safety of Lamps and Lamp Systems“. Please see “Applica-tion Note Eye Safety” for more details.At high currents, one should always avoid looking at the optical path through a focus-ing lens, since the limits imposed by Laser Class 1M can be exceeded.Electrical characteristics and op-eration of the OSRAM OSTAR Ob-servationIn addition to optimized optical behavior, the new thin film AlGaAs technology also exhib-its improved electrical characteristics, when compared to traditional standard chip tech-nologies. These improvements lead to a significantly reduced forward voltage. It also enables higher forward currents for a given junction temperature.A typical current-voltage characteristic is shown in Figure 7.Care should be taken to observe the limiting conditions specified in the data sheet and at higher power, sufficient cooling should be provided.The OSRAM OSTAR Observation consists of a current-driven component, in which small voltage fluctuations at the input can lead to significant changes in current for theApril 8, 2011page 6 of 10device and thus to changes in the emitted output power. When selecting or developing suitable driver circuitry, it is therefore rec-ommended that appropriate current stabili-zation should also be provided. To find a suitable component for this purpose please see the manufacturer homepages linked on .Figure 7: Current-Voltage characteristic of the OSRAM OSTAR Observation.The efficiency of the OSRAM OSTAR Ob-servation module which results from the total radiated light power Φe and the electrical power P = V f x I f , is plotted in Figure 8. It is optimal at around 100 mA and de-creases at lower and higher currents.This is especially true for pulse operation at I f >100 mA, since the average optical power does not remain constant when the current is doubled and the duty cycle is halved.Figure 8: Efficiency in relation to forward current I f ; T B = 25°C, t pulse = 100µs.Thermal ConsiderationsIn order to achieve reliability and optimal performance for IR light sources such as the OSRAM OSTAR Observation, appropriate thermal management is necessary.Basically, there are two principle limitations for the maximum allowable temperature. First of all, for the OSRAM OSTAR Observa-tion, the maximum allowable base plate temperature T B of 125°C must not be ex-ceeded. Secondly, the maximum junction temperature is specified to be 145°C. Since these temperatures are dependent on the operating current and mode of operation (constant current or pulsed mode), the maximum allowable currents listed in the data sheet specify a T B of up to 125°C for DC operation. Thus, for example, the maxi-mum allowable constant current is 1 A for a base plate temperature T B = 85°C and is 650 mA at 110°C. The permissible pulse handling diagram shows the maximum cur-rent allowed for various pulse conditions with given pulse length t p and duty cycle D.April 8, 2011page 7 of 10Exceeding the maximum junction tempera-ture of 145°C can lead to irreversible dam-age to the LED and to spontaneous failure of the device.Due to underlying physical inter-dependencies associated with the function-ing of light emitting diodes, a change in the junction temperature T J - within the allowable temperature range - has an effect on several LED parameters.As a result, the forward voltage, radiant flux, wavelength and lifetime of LEDs are influ-enced by the junction temperature.Influence on forward voltage V f and optical power ΦeFor LEDs, an increase in junction tempera-ture leads to both a reduction of forward voltage V F (Figure 9), and a decrease in optical power Φe (Figure 10). The resulting changes are reversible. That is, the original default values return when the temperature change is reversed.For the application, this means that the lower the temperature of the semiconductor, the higher the light output will be.Influence on reliability and lifetimeIn general, with respect to aging, reliability and performance, continually driving the LEDs at their maximum allowable junction temperature is not recommended, since with an increase in temperature, a reduction in lifetime can be observed.Figure 9: Typical forward voltage in rela-tion to base plate temperature T B (I f = 1 A, t p = 10 ms).Figure 10: Relative optical power in rela-tion to base plate temperature for various pulsed currents (t p = 10 ms).April 8, 2011page 8 of 10Determination of the module tem-perature with the integrated NTCA good approximation of the base plate temperature TB can be determined from the measured resistance of the NTC and the curve given in the reference table (Fig-ure 12).Depending on the operating conditions, the corresponding junction temperature will be ΔT = R thJB x P D (P D = electrical power dissi-pation) higher. With appropriate feedback circuitry, T B and thus the junction tempera-ture can be regulated.Figure 11: Cross section of the OSRAM OSTAR Observation.Design ExampleIn the following example, the thermal re-quirements of the heat sink for the OSRAM OSTAR Observation are examined. In Fig-ure 13, an equivalent circuit for the different thermal resistances of the module is shown. Additional information is contained in the application note "Thermal Management of OSTAR-Projection Light Source".As a starting point for the thermal evaluation, an OSRAM OSTAR Observation module (10 Chips) is driven at an operating current of I f = 1000 mA and a maximum ambient tem-perature of T A = 50°C .From the given data and information from the data sheet, the requirements for the necessary cooling can be found by means ofthe following formula:Figure 12: Typical thermistor characteris-tics for the OSRAM OSTAR Observation (NTC EPCOS 8502).Where][][][][,)()(A I V V W P T T T K T f f Module D Safety mbient A unction J ⋅≈Δ−−=ΔWithT J(unction) = Max. Junction temperature (from data sheet: T J = 145°C)T B(aseplate) = Base plate temperatureT A(mbient) = Ambient temperature (T A = 50°C)ΔT Safety = Safety temperature range (typ.10 – 20K)V f = Forward voltage (from data sheet: V f = 15.5V)I f = Forward current (I f = 1A) Æ typ. P D, Module = 15.5 WApril 8, 2011page 9 of 10ΔT = Temperature change due to P D,ModuleR th,Interface = Thermal resistance of the transition mate-rial between the OSRAM OSTAR base plate and the cooler/heat sink (e.g. thermally conductive paste ≈ 0.1 K/W)R th,JB = Thermal resistance of the OSRAM OSTAR Observation (from data sheet: R th,JB = 2.8 K/W)R th,Heat sink = Thermal resistance of the cooler/heat sink to the environmentthe thermal resistances.In this example, the maximum thermal resis-tance required for cooling of the module can be found by:With the calculated thermal resistance value at hand, a corresponding heat sink can beselected from a manufacturer (see ). Using this setup at the given operating conditions the junction temperature of the module will be at 135°C. If a lower T J is desired, the safety temperature ΔT Safety has to be increased accordingly.In addition to a thermal evaluation by means of a simulation or a computed estimate, it is generally recommended to verify and safe-guard the design with a prototype and ther-mal measurements.ConclusionDeveloped for high power operation with pulsed currents of up to five Amperes, the OSRAM OSTAR Observation IR light source achieves a light output of several Watts, depending on operating parameters.Due to operation at high power levels, ap-propriate thermal management is particularly necessary in order to dissipate the accumu-lated heat and to assure the optimal per-formance and reliability of the module.When developing applications based on the OSRAM OSTAR Observation, it is generally recommended that in addition to thermal simulations, the design should be verified and safeguarded by means of a prototype and thermal measurements.April 8, 2011page 10 of 10Don't forget: LED Light for you is your place to be whenever you are looking for information or worldwide partners for your LED Lighting project.Author: Dr. Claus Jäger, Andreas StichABOUT OSRAM OPTO SEMICONDUCTORSOSRAM is part of the Industry sector of Siemens and one of the two leading lighting manufactur-ers in the world. Its subsidiary, OSRAM Opto Semiconductors GmbH in Regensburg (Germany), offers its customers solutions based on semiconductor technology for lighting, sensor and visu-alization applications. OSRAM Opto Semiconductors has production sites in Regensburg (Ger-many) and Penang (Malaysia). Its headquarters for North America is in Sunnyvale (USA), and for Asia in Hong Kong. OSRAM Opto Semiconductors also has sales offices throughout the world. For more information go to .All information contained in this document has been checked with the greatest care. OSRAM Opto Semiconductors GmbH can however, not be made liable for any damage that occurs in connection with the use of these contents.。

Star Formation in Clusters

a rXiv:as tr o-ph/4821v111Aug24Star Formation in Clusters By S ØR E N S.L A R S E N ESO /ST-ECF,Karl-Schwarzschild Strasse 2,D-85748Garching bei M¨u nchen,Germany The Hubble Space Telescope is very well tailored for observations of extragalactic star clusters.One obvious reason is HST’s ability to recognize clusters as extended objects and measure sizes out to distances of several Mpc.Equally important is the wavelength range oﬀered by the instruments on board HST,in particular the blue and near-UV coverage which is essential for age-dating young clusters.HST observations have helped establish the ubiquity of young massive clusters (YMCs)in a wide variety of star-forming environments,ranging from dwarf galaxies and spiral disks to nuclear starbursts and mergers.These YMCs have masses and structural properties similar to those of old globular clusters in the Milky Way and elsewhere,and the two may be closely related.Several lines of evidence suggest that a large fraction of all stars are born in clusters,but most clusters disrupt rapidly and the stars disperse to become part of the ﬁeld population.In most cases studied to date the luminosity functions of young cluster systems are well ﬁt by power-laws dN (L )/dL ∝L −αwith α≈2,and the luminosity of the brightest cluster can (with few exceptions)be predicted from simple sampling statistics.Mass functions have only been constrained in a few cases,but appear to be well approximated by similar power-laws.The absence of any characteristic mass scale for cluster formation suggests that star clusters of all masses form by the same basic process,without any need to invoke special mechanisms for the formation of “massive”clusters.It is possible,however,that special conditions can lead to the formation of a few YMCs in some dwarfs where the mass function is discontinuous.Further studies of mass functions for star clusters of diﬀerent ages may help test the theoretical prediction that the power-law mass distribution observed in young cluster systems can evolve towards the approximately log-normal distribution seen in old globular cluster systems.rsen:Star Formation in Clustersexcellent blue and UV quantum eﬃciency have become available,a better understanding of“dome”seeing has led to improved image quality,and8–10m ground-based telescopes have made it possible to obtain high-quality spectra of faint objects detected in HST images.2.Why star clusters?While star clusters have been the subject of substantial interest for many years,it may be worth recalling some of the main motivations for studying them.First,there are a number of problems which make clusters interesting in their own right. These involve both their formation,subsequent dynamical evolution and ultimate fate. Atﬁrst glance,clusters appear deceptively simple:they are aggregations of a few hundred to about a million individual stars,generally constituting a gravitationally bound system (although the latter may not be true for some of the youngest systems).Yet,constructing realistic models of their structure and dynamical evolution has proven to be a major challenge,and it is only now becoming possible to carry out reasonably realistic N-body simulations including the eﬀects of stellar evolution,external gravitationalﬁelds,and the rapidly varying gravitational potential in the early phases of cluster evolution during which gas is expelled from the system(Joshi et al.2000;Giersz2001;Kroupa&Boily 2002).The models must be tested observationally,and HST data currently represent the only way to reliably measure structural parameters for extragalactic star clusters. Second,there is growing evidence that a signiﬁcant fraction of all stars form within clusters,although only a small fraction of these stars eventually end up in bound clusters (Lada&Lada2003;Fall2004).Therefore,the problem of understanding star formation is intimately linked to that of understanding cluster formation,and a theory of one cannot be complete without the other.It is of interest to investigate how the properties of star clusters might depend on environment,as this might provide important clues to any diﬀerences in the star formation process itself.In particular,HST has made important contributions towards establishing the presence of“young massive clusters”(YMCs†)in a variety of environments,which appear very similar to young versions of the old globular clusters which are ubiquitous around all major galaxies.Globular cluster formation was once thought to be uniquely related to the physics of the early Universe(e.g.Peebles& Dicke1968;Fall&Rees1985)but it now seems to be an ongoing process which can be observed even at the present epoch.Third,star clusters are potentially very useful as tracers of the stellar populations in their host galaxies.Clusters can be identiﬁed and studied at much greater distances than individual stars.In most cases,they are composed of stars which,to a very good approximation,formed at the same time and have the same metallicity.This is in contrast to the integrated light from the galaxies,which may originate from an unknown mix of stellar populations with diﬀerent ages and metallicities.Although the eﬀects of stellar evolution alone cause a cluster to fade by5-6magnitudes(in V-band)over10Gyrs (Bruzual&Charlot2003),it is in principle possible to detect clusters which have formed during the entire lifetime of galaxies,out to distances of several Mpc.In particular, globular clusters have been used extensively in attempts to constrain the star formation histories of early-type galaxies.†Also known as“super star clusters”,“populous clusters”or“young globular clusters”.rsen:Star Formation in Clusters3 3.HST and Extragalactic Star ClustersHST is almost ideally tailored for studies of extragalactic star clusters.Three main reasons for this are:•Angular resolution:clusters typically have half-light radii of2–4pc(see Section5), and can thus be recognized as extended objects out to distances of10–20Mpc with the ∼0′′.05resolution oﬀered by WFPC2or ACS.With careful modeling of the point spread function(PSF)this limit may be pushed even further.•Field size:At10Mpc,the ACS200′′×200′′ﬁeld-of-view corresponds to about10 kpc×10kpc,making it possible to cover a signiﬁcant fraction of a typical galaxy in a single pointing.•Spectral range:For studies of young stellar populations,optical and near-UV spectral coverage is essential,as discussed below.There is currently no alternative to HST on the horizon which oﬀers a similar combi-nation of capabilities.The James Webb Space Telescope(JWST),while oﬀering vastly improved eﬃciency in the IR,will oﬀer no signiﬁcant gain in resolution over HST,and will be limited to longer wavelengths.Ground-based adaptive optics(AO)can provide similar,or even better resolution than HST,but only within a small(∼20′′)isoplanatic ﬁeld of view.Furthermore,AO lacks the stable PSF of HST which is critical for many purposes(e.g.when measuring structural parameters for star clusters at the limit of the resolution),and is in any case limited to the IR(at least for now).The GALEX mission oﬀers wide-ﬁeld UV imaging,but with a spatial resolution that is inferior by far to that of HST(about5′′).The need for optical and near-UV imaging in particular deserves some additional com-ments.Figure1shows simple stellar population(SSP)model calculations(Bruzual& Charlot2003)for the evolution of the U−B,B−V and V−K broad-band colors of a single-burst stellar population.The models are shown for metallicities Z=0.02(Solar) and Z=0.004between ages of106years and1010years.As seen from theﬁgure,the U−B color is an excellent age indicator in the range from107to a few times108years, increasing by more than0.5mag and with little metallicity dependence over this age range.The B−V color,in contrast,remains nearly constant over the same age range, and oﬀers little leverage for age determinations.In practice,there are complicating prob-lems such as dust extinction,which in general will make it diﬃcult to obtain accurate age estimates from a single ing a combination of two colors,such as U−B and B−V,will make it possible to constrain both age and reddening,while at the same time being relatively insensitive to metallicity eﬀects.The relation between age and location of a cluster in the(U−B,B−V)two-color plane has been calibrated with clusters in the Large Magellanic Cloud through the so-called‘S’-sequence(Elson&Fall1985;Girardi et al.1995).For ages younger than about107years,line emission becomes important(An-ders&Fritze-v.Alvensleben2003),while the age-metallicity degeneracy(Worthey1994) becomes a diﬃculty at older ages.A more recent discussion of photometric age indica-tors,with emphasis on the importance of blue and UV data,is in Anders et al.(2004b). Use of e.g.the V−K color can help put further constraints on the metallicity and may also help constrain the ages of stellar populations in the range∼200Myr to∼500Myr (Maraston et al.2002),although the models are more uncertain and depend strongly on the stellar evolutionary tracks used in the construction of the SSP models(Girardi2000). High-resolution,wide-ﬁeld imaging in the blue and/or UV will be especially important for attempts to constrain not only the luminosity function,but also the mass function of clusters.For a long time,WFPC2was the“workhorse”on HST,and it remains the only wide-ﬁeld imager on board HST with U-band imaging capability.However,the sensitivityrsen:Star Formation in ClustersFigure1.Evolution of broad-band colors as a function of age and metallicity according to Bruzual&Charlot(2003)simple stellar population models.of WFPC2in the U-band is rather low and the detectors are steadily degrading.The Wide Field Camera3,with its panchromatic coverage,would be an ideally suited instrument for such studies.4.Setting the stage:early developmentsEven within the Local Group,it has long been known that the traditional distinction between open and globular clusters that can be applied fairly easily in the Milky Way breaks down in some other galaxies.The classical example is the“blue globular”clusters in the Large Magellanic Cloud,which are not easily classiﬁed as either open or globular clusters.The most massive of these objects have masses up to∼105M⊙(Elson&Fall 1985;Fischer et al.1992;Richtler1993;Hunter et al.2003)similar to the median mass of old globular clusters and about an order of magnitude more massive than any young open cluster known in the Milky Way.Yet,these objects have young ages,and are still being produced today by the LMC.Similar clusters have been found in M33(Christian &Schommer1982,1988).A good indication of the status of research in extragalactic young star clusters shortly prior to HST is provided by Kennicutt&Chu(1988;hereafter KC88).These authors compiled observations of what they refer to as“young populous clusters”(PCs)in14 galaxies for which data was available at that time.Half of the galaxies studied by KC88 were Local Group members(Milky Way,LMC,SMC,M33,M31,NGC6822and IC1613). As noted by KC88,a severe diﬃculty in comparing observations of PCs in diﬀerent galaxies,made by diﬀerent authors,is the widely variable completeness of the surveys, and the diﬀerent deﬁnitions of such clusters.KC88adopted a(somewhat arbitrary) deﬁnition of a young PC as an object with an estimated mass>104M⊙and a colorrsen:Star Formation in Clusters5 B−V<0.5.They noted a conspicuous deﬁciency of populous clusters in the Milky Way and M31,the two only large Sb/Sbc-type spirals in the sample,and suggested that this might be linked to the deﬁciency of giant H ii regions in the same two galaxies.By comparing the relative numbers of PCs and giant H ii regions in their sample of galaxies, KC88concluded that PCs may indeed form inside such regions,but not all giant H ii regions produce bound clusters.This is very much in line with recent indications that only a small fraction of star clusters of any mass remain bound(Fall2004).The galaxies which did contain PCs were all late-type,though not all late-type galaxies were found to contain PCs.A signiﬁcant exception is the Local Group dwarf irregular IC1613,which contains few if any star clusters at all(van den Bergh1979;Hodge1980)in spite of some on-going star formation.The near-absence of star clusters in IC1613may be as important a clue to the nature of the cluster formation process as the abundant cluster systems in starbursts and merger galaxies(Section5).To a large extent,research in old globular clusters(GCs)remained detached from that of YMCs until fairly recently.It was well-known that early-type galaxies typically have many more GCs per unit host galaxy luminosity(Harris&van den Bergh1981)than spirals and irregulars,a fact that was recognized as a problem for the idea that early-type galaxies form by mergers of gas-rich spirals(van den Bergh1982).Schweizer(1987) proposed that this problem might be solved if new GCs form during the merger.This idea was further explored by Ashman&Zepf(1992),who predicted that the resulting merger product should contain two distinct GC populations:one metal-poor population inherited from the progenitor galaxies,and a new metal-rich population formed in the merger.The two GC populations should be identiﬁable in the color distributions of the resulting GC systems.Two highly inﬂuential discoveries soon followed:Bimodal color distributions were discovered in several GC systems around early-type galaxies(Zepf &Ashman1993;Secker et al.1995;Whitmore et al.1995),and highly luminous,com-pact young star clusters were found in ongoing or recent mergers like the Antennae and NGC7252(Whitmore et al.1993;Whitmore&Schweizer1995).In retrospect,it had al-ready been known for a long time that even the metallicity distribution of the Milky Way GC system is strongly bimodal(Zinn1985).The mean metallicities of the two modes in the Milky Way are in fact quite similar to those seen in early-type galaxies.The Milky Way is unlikely to be the result of a major merger,and there are also other indications that not all properties of GC systems in early-type galaxies can be explained by a naive application of the merger model.Alternative scenarios have later been put forward to ex-plain the presence of multiple GC populations(e.g.Forbes et al.1997;Cˆo t´e et al.1998), but it is beyond the scope of this paper to discuss any of these in prehensive discussions can be found e.g.in Harris(2001)and Kissler-Patig(2000).Nevertheless,the discovery of young globular cluster-like objects in ongoing mergers was a tantalizing hint that it might be possible to study the process of globular cluster formation close-up at the present epoch,and not just from the fossil record.5.Extragalactic Star Clusters in Diﬀerent EnvironmentsTables1–5are an attempt to collect a reasonably complete list of galaxies where YMCs have been identiﬁed(until∼May2004),along with some pertinent references.For each galaxy,the main facilities used for the observations are listed,although in many cases it is impossible to give a comprehensive listing.Standard abbreviations(ACS,FOC,GHRS, STIS,NICMOS,WFPC,WFPC2)are used for HST instruments.Other abbreviations are WIYN(Wisconsin Indiana Yale NOAO3.5m),UKIRT(United Kingdom Infra-Red Telescope),NTT(ESO3.5m New Technology Telescope),CFHT(3.6m Canada-Francersen:Star Formation in ClustersHawaii Telescope),DK154(Danish1.54m at ESO,La Silla)and NOT(2.56m Nordic Optical Telescope).The level of detail provided in diﬀerent studies varies enormously–in some cases,identiﬁcations of YMCs are only a byproduct of more general investigations of galaxy properties(e.g.Meurer et al.1995;Scoville et al.2000)while other studies are dedicated analyses of cluster systems in individual galaxies.Galaxies marked with an asterisk(⋆)are used later(Section6.2)when discussing luminosity functions.In the following I brieﬂy discuss a few illustrative cases from each table and then move on to discuss more general properties of young cluster systems.5.1.Starburst galaxiesThe richest populations of YMCs are often found in major mergers(Section5.2).How-ever,there are also examples of YMCs in starbursts which are not directly associated with major mergers,although they may in some cases be stimulated by more benign interactions or accretion of companion satellites(Table1).In the case of NGC7673, for example,Homeier&Gallagher(1999)argue that the morphological features of the galaxy point toward a minor merger,while the starburst in M82may have been triggered by tidal interactions with M81.M82is also noteworthy for being theﬁrst galaxy in which the term‘super star cluster’was used.It was introduced by van den Bergh(1971), who was careful to point out that the nomenclature was not intended to imply that these objects are necessarily bound.The presence of SSCs in M82was conﬁrmed by O’Connell et al.(1995)who identiﬁed about100clusters in WFPC images.An example of a starburst which is unlikely to be triggered by an interaction is NGC5253,which is located about600kpc from its nearest neighbor,M83(Harris et al.2004).One of theﬁrst surveys to provide a systematic census of star clusters in a sample of starburst galaxies was the work by Meurer et al.(1995),who observed9galaxies with HST’s Faint Object Camera(FOC).Meurer et al.noted that a high fraction,on average about20%,of the UV luminosity in these starbursts originated from clusters or compact objects,and a hint of a trend for this fraction to increase with the underlyingrsen:Star Formation in Clusters7UV surface brightness.They also measured cluster sizes similar to those of Galactic globular clusters,and found the luminosity functions to be well represented by a power-law dN(L)/dL∝L−αwithα≈2.YMCs have been identiﬁed in several nuclear and circumnuclear starburst regions,often associated with barred spiral galaxies(Table2).Maoz et al.(1996)studied5circumnu-clear starbursts and found that as much as30%–50%of the UV light came from compact, young star clusters with half-light radii<5pc and estimated masses up to about105M⊙. Again they found the luminosity functions to be well approximated by power-laws with slopeα≈2.Buta et al.(2000)found a much steeper slope(α=3.7±0.1)in their study of the circumnuclear starburst in NGC1326,but noted that their sample might be contaminated by individual supergiant stars.In some cases the ring-like structure of the nuclear starburst is not quite so evident.Watson et al.(1996)discovered4lu-minous clusters in the central starburst region of NGC253,the brightest of which has M V=−15,an inferred mass in excess of1.5×106M⊙and a half-light radius of2.5pc. However,these clusters may be part of a compact ring-like structure with a radius of about50pc(Forbes et al.2000).Most of the clusters in the nuclear starburst of M83 are also located within a semicircular annulus(Harris et al.2001),but again the ring is more poorly deﬁned.5.2.MergersMany of the most spectacular YMC populations have been found in merger galaxies. NGC1275was one of theﬁrst galaxies in which HST data conﬁrmed the existence of YMCs,although at least one object in this galaxy was already suspected to be a massive cluster based on ground-based data(Shields&Filippenko1990).With the Planetary Camera on HST,Holtzman et al.(1992)identiﬁed about60cluster candidates with ab-solute magnitudes up to M V=−ing WFPC2data,Carlson et al.(1998)identiﬁed about3000clusters,of which about1200have blue integrated colors and estimated ages between0.1and1Gyr.The young clusters had estimated masses and sizes similar to those of old globular clusters,although Brodie et al.(1998)found that the Balmer line equivalent widths measured on spectra of5clusters were too strong to be consistent with standard SSP models,unless a stellar mass function truncated at2M⊙−3M⊙wasrsen:Star Formation in Clustersadopted.With accurate modeling of the HST point spread function and high dispersion spectroscopy with8–10m class telescopes,it might be possible to constrain the virial masses of some of the brightest clusters,and thereby provide independent constraints on their stellar IMF.While NGC1275may have experienced a recent merger/accretion event(Holtzman et al.1992),it is hardly one of the classical“Toomre”mergers(Toomre&Toomre 1972).One of the nearest ongoing,major mergers is the“Antennae”NGC4038/39, where HST observations have revealed a rich population of luminous,compact young star clusters with typical half-light radii∼4pc(Whitmore&Schweizer1995;Whitmore et al.1999).The brightest of them reach M V≈−14and have estimated masses close to 106M⊙(Zhang&Fall1999).Similar rich populations of YMCs have been found in many other mergers,like NGC3256where Zepf et al.(1999)identiﬁed about1000compact bright,blue objects on WFPC2images within the central7kpc×7kpc region.Again, the young clusters contribute a very signiﬁcant fraction(15%–20%)of the blue light within the starburst region.Zepf et al.(1999)estimated half-light radii of5–10pc for the clusters in NGC3256,somewhat larger than for the Antennae,but note that1PC pixel corresponds to a linear scale of8pc at the distance of NGC3256,so that thersen:Star Formation in Clusters9clusters are only marginally resolved.Interestingly,only a shallow trend of cluster sizeversus luminosity was found,with radius r scaling with luminosity L roughly as r∝L0.07.NGC7252is a somewhat more advanced system than NGC3256or the Antennae.Miller et al.(1997)date the cluster system at between650Myr and750Myr.Remark-ably,both photometry and dynamical measurements yield a mass of about8×107M⊙for the most massive object(W3)(Maraston et al.2004),making it about an order of magnitude more massive than any old globular cluster in the Milky Way.With a half-light radius of17.5±1.8pc,this object is much larger than a normal star cluster,andmay be more closely associated with the“Ultra Compact Dwarf Galaxies”in Fornax (Hilker et al.1999;Drinkwater et al.2003).5.3.Dwarf/Irregular galaxiesThe bright“central condensations”in NGC1569were noted already by Mayall(1935)onplates taken with the36inch Crossley reﬂector at Lick Observatory,though Arp&Sandage(1985) were probably theﬁrst to recognize them as likely star clusters.At a distance of only∼2Mpc(Makarova&Karachentsev2003),these clusters appear well resolved on HST images with half-light radii of about2pc(O’Connell et al.1994;de Marchi et al.1997).One of the clusters,NGC1569-A,is actually a double cluster,and STIS spectroscopyhas shown that one component exhibits Wolf-Rayet features while the other componentis devoid of such features,suggesting an age diﬀerence of a few Myrs between the two components(Maoz et al.2001b).Using high-dispersion spectroscopy from the NIRSPEC spectrograph on the Keck II telescope,Gilbert&Graham(2003)derived dynamical mass estimates of about0.3×106M⊙for each of the two components of NGC1569-A,and0.18×106M⊙for NGC1569-B,again very similar to the typical masses of old globular clusters,and consistent with the clusters having“normal”stellar mass functions(see also Section6.4).A peculiar feature of the NGC1569cluster population is that the next brightest clustersafter NGC1569-A and NGC1569-B are more than2magnitudes fainter(O’Connell etrsen:Star Formation in Clustersal.1994).An even more dramatic discontinuity in the luminosity function is seen in NGC1705which has only a single bright cluster,and in NGC4214there is a gap of about1.5mag from the brightest2clusters down to number3(Billett et al.2002). Interestingly,while the clusters in NGC1569and NGC1705are young(∼107years),the two clusters in NGC4214are both about250Myrs old(Billett et al.2002),demonstrating that massive clusters are capable of surviving for substantial amounts of time at least in some dwarf galaxies.5.4.Spiral galaxy disksMost of the YMCs discussed in the preceding sections are located in environments that are peculiar in some way,or at least diﬀerent from what we see in the solar neighborhood. Thus,it is tempting to speculate that the absence of YMCs in the Milky Way indicates that their formation somehow requires special conditions.There is,however,increasing evidence that YMCs can form even in the disks of spiral galaxies.Table5lists a number of nearby spirals in which YMCs have been identiﬁed.A few(e.g.M51)are clearly involved in interactions,but none of them are disturbed to a degree where they are not clearly recognizable as spirals.The nuclear starburst in M83was already mentioned in Section5.1,but there is also a rich population of young star clusters throughout the disk (Bohlin et al.1990;Larsen&Richtler1999),the most massive of which have masses of several times105M⊙.An even more extreme cluster is in NGC6946,with a dynamical mass estimate of about1.7×106M⊙(Larsen et al.2001).The disks of spiral galaxies can evidently form star clusters with masses as high as those observed in any other environment,including merger galaxies like the Antennae and starbursts like M82. Most of the spirals in Table5are type Sb or later,but one exception is NGC3081.In this barred S0/Sa-type spiral,Buta et al.(2004)detected a number of luminous young clusters in the inner Lindblad resonance ring at5kpc.Buta et al.(2004)found rather large sizes for these clusters,with estimated half-light radii of about11pc.This is much larger than the typical sizes of Milky Way open and globular clusters and indeed of YMCs found in most other places,and raises the question whether these objects might be related to the“faint fuzzy”star clusters which are located in an annulus of similar radius in the lenticular galaxy NGC1023(Larsen&Brodie2000;Brodie&Larsen2002), but have globular cluster-like ages.rsen:Star Formation in Clusters11 6.General properties of cluster systemsJust how similar are the properties of star clusters in diﬀerent environments,and what might they tell us about the star formation process?Objects like NGC1569-A appear extreme compared to Milky Way open clusters or even to young LMC clusters: O’Connell et al.(1994)estimate that NGC1569-A has a half-light surface brightness over 65times higher than the R136cluster in the LMC,and1200times higher than the mean rich LMC cluster after allowing for evolutionary fading.Do such extreme objects constitute an altogether separate mode of star/cluster formation,or do they simply represent a tail of a distribution,extending down to the open clusters that we encounter locally?And are YMCs really young analogs of the old GCs observed in the Milky Way and virtually all other major galaxies?6.1.Luminosity-and mass functionsOne of the best tools to address these questions is the cluster mass function(MF).In the Milky Way and the Magellanic Clouds,the MF of young star clusters is well approximated by a power-law dN(M)/dM∝M−αwhereα≈2(Elmegreen&Efremov1997;Hunter et al.2003).This is deceptively similar to the luminosity functions derived in many young cluster systems,but it is important to recognize that luminosity functions are not necessarily identical,or even similar to the underlying MFs(unless the age distribution is a delta function).Unfortunately,MFs are diﬃcult to measure directly.The only practical way to obtain mass estimates for large samples of clusters is from photometry,but because the mass-to-light ratios are strongly age-dependent masses cannot be estimated without reliable age information for each individual cluster.As discussed in Section3,this is best done by including U-band imaging,which is costly to obtain in terms of observing time. So far,MFs have only been constrained for a few,well-studied systems.In the Antennae, Zhang&Fall(1999)found a power-law shape with exponentα≈2over the mass range 104M⊙to106M⊙,similar to the MF of young LMC clusters.Bik et al.(2003)ﬁnd α=2.1±0.3over the range103M⊙to105M⊙for M51,and de Grijs et al.(2003a)ﬁnd α=2.04±0.23andα=1.96±0.15in NGC3310and NGC6745.The many studies which have found similar power-law luminosity functions are of course consistent with these results,but should not be taken as proof that the MF is as universal as the LF.Conversely,any diﬀerences in the LFs observed in diﬀer-ent systems would not necessarily imply that the MFs are diﬀerent.There are some hints that slight LF variations may be present:Elmegreen et al.(2001)ﬁnd LF slopes ofα=1.58±0.12andα=1.85±0.05in NGC2207and IC2163,while Larsen(2002) and Dolphin&Kennicutt(2002)ﬁnd somewhat steeper slopes(α=2.0−2.5)in several nearby spiral galaxies.While Whitmore et al.(1999)ﬁndα=2.12±0.04for the full sample of Antennae clusters,there is some evidence for a steepening at brighter mag-nitudes withα=2.6±0.2brighter than M V=−10.4.However,measurements of LF slopes are subject to many uncertainties,as completeness and contamination eﬀects can be diﬃcult to fully control,and it is not presently clear how signiﬁcant these diﬀerences are.More data is needed.Another important question is how the MF evolves over time.While the evidence avail-able so far indicates that the MF in most young cluster systems is well approximated by a uniform power-law with slopeα≈2down to the detection limit,old GC systems show a quite diﬀerent behavior.Here,the luminosity function is wellﬁt by a roughly log-normal distribution with a peak at M V∼−7.3(about105M⊙for an age of10–15 Gyr)and dispersion∼1.2mag(e.g.Harris&van den Bergh1981).Thus,old globu-lar clusters appear to have a characteristic mass of about∼105M⊙,while there is no characteristic mass for young clusters.This diﬀerence might seem to imply fundamen-。

2023-2024学年全国天域名校协作体高三下学期3月联考英语试题

2023-2024学年全国天域名校协作体高三下学期3月联考英语试题1. Why is the man phoning the woman?A．To ask her out for dinner.B．To tell her he will be late.C．To inform her he has been in an accident.2. What does the man like his toast with?A．Much butter. B．Less butter. C．Nothing on it.3. What is the man going to buy his father?A．A new digital camera. B．An antique desk. C．An old film camera.4. How fast did the man drive on the straightaway?A．180 km/h. B．200 km/h. C．220 km/h.5. How is the man feeling?A．Sad. B．Awkward. C．Proud.听下面一段较长对话，回答以下小题。

6. When was the last time the speakers were in a theater?A．January. B．March. C．July.7. What stage show did the speakers watch on television?A．Cats. B．Romeo and Juliet. C．Hamilton.听下面一段较长对话，回答以下小题。

8. Where does the conversation probably take place?A．In a bar. B．In a hotel. C．In a restaurant.9. How much is the barbecue?A．$10. B．$20. C．$30.10. What is the woman going to do first?A．Have a drink. B．Take a shower. C．Eat some food.听下面一段较长对话，回答以下小题。

Unihedron SQM-L 天空光强计说明书

Thank you for purchasing a Sky Quality Meter (SQM-L) from Unihedron! FeaturesThe SQM-L has the following features:It is sensitive only to visual light (there is a near-infrared blocking filter in front of the sensor).The effects of temperature on the “dark frequency" of the sensor are removed.The effects of temperature on the microcontroller oscillator are removed. It is protected against accidental reversal of battery polarity.Each SQM-L is calibrated using a NIST-traceable light meter.The absolute precision of each meter is believed to be±10%(±0.10 mag/arcsec²).The difference in zeropoint between each calibrated SQM-L is typically±10%(±0.10 mag/sq arcsec)The brightness of the numeric LEDdisplay has two(automatic)settings.Under dark skies, you won't have yourdark adaption ruined by use of yourSQM-L! Under urban skies, the displaywill be correspondingly brighter.A repeating audible beep indicates whena measurement is in progress.Any kind of 9V battery is usable. TheSQM-L contains a voltage regulator topower the sensor,microcontroller andother components.After reading is taken and displayed, themeter automatically turns itself off.The Half Width Half Maximum(HWHM) of the angular sensitivity is10∼°. The Full Width Half Maximum(FWHM) is then 20∼°. The sensitivityto a point source 19°off-axis is a∼factor of 10 lower than on-axis. A pointsource 20° and 40° off-axis would∼∼register 3.0 and 5.0 magnitudes fainter,respectively.✳ ✵ ✳Quick StartThe SQM-L is very simple to use. Pointthe lens towards the zenith. Press the Startbutton once and release.Under urbanskies, a reading will be displayed almostimmediately. Under the very darkestconditions (no moon in the sky, far fromcivilization)the meter may take up to aminute to complete its measurement.Please ensure that you maintain theorientation of the meter until the reading isdisplayed.The SQM-L's reading is indicative of thesky brightness within its field of view.There must be no direct illumination orshading of the sensor by a terrestrial lightsource if the reading is to be meaningful.✳ ✵ ✳Typical ReadingsMagnitudes per square arcsecond is alogarithmic measurement. Therefore largechanges in sky brightness correspond torelatively small numerical changes. Adifference of 1 magnitude is defined to be afactor of(100)(1/5)in received photons.Therefore a sky brightness 5.0 mag/arcsec²fainter corresponds to a reduction inphoton arrival rate of a factor of 100.The following schematic gives a roughidea of of how to interpret the readings:At the darkest sites, natural variations inconditions such as airglow and thebrightness of the zodiacal light are limitingfactors.✳ ✵ ✳Temperature readingThe temperature in °C then °F aredisplayed when you press and hold thebutton a second time. Also, the model andserial number are displayed after thetemperature.✳ ✵ ✳Care of your SQM-LThe SQM-L is a fairly simple and robustdevice.Avoid dropping,immersing,andcompressing it and it will give you years ofdependable service.Keep the faceplateclean and ensure that the battery still hasuseful capacity. If you have left your SQM-L for a long period of time (i.e. years) andsee a white, powdery substance around oneof the battery contacts,your battery willneed to be replaced and the contactscleaned before you can expect reliableoperation.The SQM-L should not be negativelyaffected by dew during normal operationEXCEPT for the reduction in receivedlight by the sensor.Make sure that thesensor faceplate has been wiped beforemaking measurements.During storage, make sure that the push-button is not being continuously pressedsince the meter will draw current from thebattery and drain it in that situation.Do not point the meter at the Sun.✳ ✵ ✳TroubleshootingAfter I push the button,no reading is displayed.Are you in a very dark location?Yes → The Sky Quality Meter may take up to a minute to acquire a reading when the sky is very dark. If your meter is operating properly, you will here a soft beeping sound while the measurement is in progress.When complete,the sky brightness will be displayed for a fixed number of seconds.No → Your 9V battery may need to be replaced.ORThe connector to your 9V battery may be loose.If, after you have checked for both of these possibilities and your SQM-L still won't display a reading under normal operating conditions, contact Unihedron for further information and a possible replacement.I don't know how to make sure the SQM-L is off.The SQM-L functions in such a way that it is only temporarily on and turns itself off automatically. This is a design feature to maximize battery life.The readings don't repeat exactly.Are you pointing the SQM-L in thesame direction each time? Under darkconditions, you must keep the SQM-Lpointed in the same direction until thereading appears on the LED display.Your SQM-L must be pointed at anangle sufficiently high above thehorizon that it will not detect lightdirectly from terrestrial sources(cars,buildings,streetlights).It is normallythe zenith sky brightness which ismeasured.The readings do not change when pointingto various parts of the night sky.Each SQM-L reading must be initiatedby pressing the button.The displayedreading will stay on for10secondsbefore shutting down. After the unit hasshut down, press the button to initiateanother reading.The readings are numerically lower(brighter) than expected.Make sure that no stray light from streetlights or other sources directlyilluminates the lens/sensor.The readings are numerically higher(darker) than expected.Make sure that nothing shades the fieldof view of the lens/sensor (such as a tallstand of trees or the side of a building).When I use the meter during the day, all Isee is a on the display.The SQM-L has a fantastically largerange over which it will report accuratesky brightnesses.However,to besensitive in the darkest conditions, it isnecessary to sacrifice the ability torecord daytime sky surface brightnesses.Normal lux meters can be used in suchcircumstances once the effective solidangle for the lux meter's sensor isknown.Theindicates that thesensor is saturated.All I see is a on the display.Theindicates that the sensorwas unable to produce a reading. Thiscan occur in a light-tight dark room or ifthe sensor is faulty.Sometimes the first reading is different.As the temperature of the unit changesslightly due to being powered up, thevery first reading may be slightly higherthan the following readings. Ignore thisfirst reading and average the followingones for the most accurate value.Other scalesTo convert the SQM-L mag/arcsec²reading to cd/m², use the followingformula:[cd/m2] = 10.8×104 × 10(0.4*[mag/arcsec2])Unanswered QuestionsHelp us to inform you and other customersbetter by forwarding unanswered questionsabout the SQM-L and measuring lightpollution to:******************Further InformationCheck the website forupdates and additional information.Mailing ListJoin the SQM mailing list for notificationsand to share experiences with other usersby sending an e-mail to:sqm -***********************✳ ✵ ✳Contact InformationUnihedron4 Lawrence AveGrimsby, ON L3M 2L9CanadaTel: (905) 945-1197✳ ✵ ✳Unihedron is a proud member of theInternational Dark-Sky Association()and supports its goals.Please consider joining to help preserve thebeauty of the night sky for futuregenerations.✳ ✵ ✳WarrantyUnihedron warrants this product 1 year.✳ ✵ ✳Last updated: February 11, 2008。

Unit+2+Grammar+and+usage+课件牛津译林版(2020)选择性必修第三册

It is not clear how long he would stay on the space station.
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4 What the lunar probe did not land was a worry for the people
back on the Earth.
主从为陈述句，不缺成分和意义
What makes the Chinese people happy is that China has successfully launched a space rocket.
1.It is a pity that we might miss the opening scene. 2.They lost their way in the forest and what made matters worse was _t_h_a_t_
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Below is a newspaper feature article on osefntteelenscceospwesittho)seuxbpjleocrtecslapuascees..Find the
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1. That Hubble is based in space allows it to see further than groundbased telescopes, … 2. Whether life on other planets does exist is yet to be proved, but the signs are promising... 3. It is clear that telescopes are crucial tools for space exploration and that developing the required technology will help astronomers all over the world make exciting discoveries. 4. How much we will learn from the telescopes is merely limited by our imagination.

The PDS starburst galaxies

The starburst nature of the PDS galaxies is established by comparing their LIR/LB ratios and IRAS colors with a sample of emission line galaxies from the literature already classiﬁed as starburst galaxies. The starburst galaxies show an excess of FIR luminosity and their IRAS colors are signiﬁcantly diﬀerent from those of Seyfert galaxies – 99% of the starburst galaxies in our sample have a spectral index α(60, 25) < −1.9. As opposed to Seyfert galaxies, very few PDS starbursts are detected in X-rays.
The PDS starburst galaxies
Roger Coziol1,2, Carlos A. O. Torres1, Germano R. Quast1, Thierry Contini3 and
Emmanuel Davoust4
Received
; accepted
arXiv:astro-ph/9807136v1 14 Jul 1998
–2–
ABSTRACT
We discuss the nature of the galaxies found in the Pico dos Dias Survey (PDS) for young stellar objects. The PDS galaxies were selected from the IRAS Point Source catalog. They have ﬂux density of moderate or high quality at 12, 25 and 60 µm and spectral indices in the ranges −3.00 ≤ α(25, 12) ≤ +0.35 and −2.50 ≤ α(60, 25) ≤ +0.85. These criteria allowed the detection of 382 galaxies, which are a mixture of starburst and Seyfert galaxies. Most of the PDS Seyferts are included in the catalog of warm IRAS sources by de Grijp et al. (1987). The remaining galaxies constitute a homogeneous sample of luminous (log(LB/L⊙) = 9.9 ± 0.4) starburst galaxies, 67% of which were not recognized as such before.

牛郎星和织女星在哪里英语作文

牛郎星和织女星在哪里英语作文Altair and Vega are known as the stars representing the legendary couple, the Cowherd and the Weaver Girl in Chinese mythology.They are visible to the naked eye in a specific part of the sky.To locate Altair and Vega, one needs to look towards the celestial sphere in the right direction and at the right time.These stars are typically visible in the summer months, adding a touch of magic to the night sky.With a clear and dark sky, it is possible to spot Altair and Vega using simple stargazing tools or even with the naked eye.They are often mentioned in astronomical literature and guides, providing specific directions on how to find them.The story of the Cowherd and the Weaver Girl adds an extra layer of charm to these stars.It makes them not just celestial bodies, but symbols of love and separation.Observing Altair and Vega can be a truly awe-inspiring experience, as it connects us to the myths and legends of ancient times.It reminds us of the wonder and beauty that lies beyond our planet.。

西宁2024年04版小学4年级第2次英语第1单元测验卷(有答案)

西宁2024年04版小学4年级英语第1单元测验卷(有答案)考试时间：100分钟（总分:140）B卷考试人：_________题号一二三四五总分得分一、综合题(共计100题)1、填空题：My _______ (狗) loves to play fetch with a stick.2、What do you use to cut paper?A. GlueB. ScissorsC. TapeD. Ruler3、听力题：A _______ is a measure of how much mass is contained in a given volume.4、填空题：My mom loves __________ (知识分享).5、听力题：Near-Earth objects (NEOs) are asteroids and comets that can come close to _______.6、听力题：Chemical reactions can release or absorb ________.7、填空题：The first female Prime Minister of the UK was ________ (玛格丽特·撒切尔).8、What is the capital of Hungary?A. BudapestB. DebrecenC. SzegedD. Miskolc答案:AA suspension is a mixture where particles are _____ in a liquid.10、填空题：We celebrate New Year in ______.11、听力题：Light travels in straight ______.12、填空题：A parrot can ________________ (说话).13、填空题：My aunt has a __________ (幽默感).14、What is the primary color of the sky?A. GreenB. BlueC. YellowD. Red15、填空题：Crickets make a _________ (声音) at night.16、填空题：My town is located near a __________ (湖).17、听力题：The chemical symbol for europium is ______.18、听力题：A __________ can often be found building dams in rivers.19、填空题：My favorite holiday is _______ because of the food.20、听力题：A ____ is known for its hopping abilities and is often found in gardens.21、填空题：The ________ is a gentle friend that listens.22、听力题：I can ___ (ride) a horse.The ancient Romans built __________ (道路) to connect their empire.24、听力题：They are _____ (playing) soccer.25、ers are known for their vivid ______ and ability to thrive in gardens. (某些花以其鲜艳的颜色和在花园中茁壮成长而闻名。

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a r X i v :a s t r o -p h /9810384v 1 23 O c t 1998to appear in Astronomical Journal (Jan 1999)Starburst or Seyfert?Using near-infrared spectroscopy tomeasure the activity in composite galaxiesTanya L.Hill School of Physics,University of Sydney Charlene A.Heisler Mount Stromlo and Siding Spring Observatories Ralph Sutherland Australian National University Astrophysical Theory Centre and Mount Stromlo and Siding Spring Observatories and Richard W.Hunstead School of Physics,University of Sydney ABSTRACT We present near-infrared spectra for a sample of galaxies with ambiguous optical emission line ratios.These galaxies fall between starbursts and Seyferts in the usual optical diagnostic diagrams.We ﬁnd a similar result with the near-infrared emission line ratios,which suggests that the galaxies are composite,containing both a starburst and AGN component.Furthermore,CO absorption,produced in late-type stars,is detected within the sample,although at a weaker level than is typical for starburst galaxies.We conclude that the CO feature isbeing diluted by a contribution from an AGN,thereby conﬁrming the composite nature of these galaxies.Subject headings:galaxies:active —galaxies:Seyfert —galaxies:starburst1.IntroductionThe mechanisms by which emission lines are excited in starburst and Seyfert galaxies are not fully understood.In simplistic terms,the gas in starburst galaxies is photoionised by young,hot OB stars,whereas in Seyfert galaxies,a class of Active Galactic Nuclei (AGN),the energy is believed to be derived from photoionisation of accreting material around a supermassive black hole and the ionising spectrum takes the form of a powerlaw continuum.There is current debate,however,over the contribution of star formation to the energy dynamics of Seyfert galaxies.At one extreme is the starburst model for AGN(Terlevich et al.1992),which is claimed to be able to reproduce the properties of an AGN using compact supernova remnants(SNRs)alone,eliminating the need for a black hole entirely. Furthermore,recentﬁndings(Boyle&Terlevich1998)have suggested that star formation plays an important role in QSOs(more luminous AGNs),due to the striking similarity between the evolution of the QSO luminosity density and galaxy star formation rate.It is well established that there are some galaxies in which both star formation and an AGN contribute to the observed emission,and their importance for understanding galaxy energetics is only just being realised.For example the Seyfert2galaxy,NGC1068,has an extended luminous star forming region located≈3kpc from the central AGN(Telesco &Decher1988),which provides half of the total luminosity of NGC1068(Lester et al. 1987;Telesco et al.1984).Recent observations with the MPE3D near-infrared(NIR) imaging spectrometer have further revealed a stellar core≈50pc in size that provides7% of the nuclear bolometric luminosity of NGC1068(Thatte et al.1997).Likewise,Genzel et al.(1995)have uncovered a circumnuclear ring of star formation in the Seyfert1galaxy NGC7469.In this case,almost two-thirds of the total luminosity of NGC7469arises from starburst activity.Furthermore,a compact starburst was found in Mrk477,an extremely luminous Seyfert2galaxy with a hidden Seyfert1component.The uncovered starburst is believed to have a bolometric luminosity comparable to the Seyfert component(Heckman et al.1997).Composite galaxies,i.e.,those with both a starburst and Seyfert component,have implications for evolutionary scenarios in which one type of activity may initiate the other. For example,Norman&Scoville(1988)showed theoretically that a black hole can be formed through the evolution of a massive starburst at the centre of a galaxy.In the study of NGC7469,Genzel et al.(1995)interpret a blue-shifted ridge between the nucleus and star formation ring as infalling gas providing fuel for the AGN.Alternatively,an AGN could disturb the gas in a galaxy in such a way as to trigger star formation,a mechanism proposed for the radio galaxies3C285(van Breugel&Dey1993)and Minkowski’s Object(van Breugel et al.1985).We have undertaken a search of the literature for composite galaxies based on optical spectroscopy of their nuclei as outlined in Section2.Further in Section2,we present photoionisation modelling of the optical emission line ratios that are commonly used as diagnostics.To gain insight into the sources of ionisation present within the galaxieswe begin our investigation with NIR spectroscopy.Details of the observations and data reduction are given in Section3.The obvious advantage of the NIR is that dust extinction is a factor of10lower than in the optical,thereby making it possible to probe further into the central dust-obscured regions of galaxies where the activity occurs.We explore several emission and absorption features as possible NIR ionisation diagnostics in Section4,with the help of photoionisation models.Throughout this paper,we adopt H0=75km s−1 Mpc−1and q0=0.5.2.Sample Selection2.1.Optical Emission Line DiagnosticsStarburst galaxies and narrow line AGNs both exhibit emission lines with FWHM<1000km s−1.Diﬀerences in the degree of ionisation are traditionally believed to be due to the hard ionising spectrum of the AGN compared with the UV radiation from stars.It is generally accepted that AGNs contain an extended partially ionised zone.The hard X-ray component of an AGN produces high energy X-ray and far-UV photons with long mean free paths that can extend beyond the ionised region to create a large zone of partially ionised gas.For starburst galaxies,on the other hand,the transition between neutral and fully ionised gas is sharp;the UV radiation is readily absorbed by neutral hydrogen,and hence the partially ionised zone is small.The result is that the forbidden lines,e.g.[O III],[N II], [S II]and[O I],which are collisionally excited and produced mostly in partially ionised zones,are stronger in Seyfert galaxies than in starbursts.On the other hand,hydrogen recombination lines are stronger in starburst galaxies due to the H II regions surrounding young hot OB stars.Optical emission line ratios are traditionally used to determine the photoionising source in narrow emission line galaxies.Veilleux&Osterbrock(1987)found empirically that the best optical diagnostic ratios are those of[O III]/Hβcompared with[N II]/Hα,[S II]/Hαand[O I]/Hα.These line ratios are sensitive to the presence of partially ionised gas and are comprised of easily observable lines spanning a relatively narrow wavelength range,thereby minimising reddening uncertainties.The diagnostic diagrams segregate the galaxies in the manner predicted by photoionisation models which use OB stars as the ionising source for starburst galaxies and a powerlaw input spectrum for AGN(Veilleux&Osterbrock1987).However,not all galaxies can be easily classiﬁed using the diagnostic diagrams(eg.Ashby et al.1995).In addition,photoionisation may not be the only mechanism operating to produce the narrow emission line spectra.The importance of shock excitation over photoionisation has been evaluated within the optical regime by Dopita&Sutherland(1996),who show that the spectral characteristics of many Seyfert galaxies can be produced by shock excitation alone. Shocks can have both a stellar and non-stellar origin,and are believed to be associated with supernovae in starburst galaxies and radio jets in AGN.Thus,it is important to study in detail those galaxies with ambiguous optical line ratio diagnostics in an eﬀort to understand the mechanisms involved in producing the observed emission lines.If both star formation and an AGN are powering such galaxies,these galaxies may prove to be important for understanding the relationship,if any,between the two energy sources.2.2.Sample Selection CriteriaA search of the literature was undertaken for large optical spectroscopic surveys from which to select composite galaxy candidates,that is,galaxies that potentially contain both intense star formation and an AGN.We expect that such galaxies would have ambiguous optical line ratios(intermediate to the starburst and AGN classes)and therefore used the optical diagnostic diagrams to choose galaxies that(1)fell within±0.15dex of the boundary line determined by Veilleux&Osterbrock(1987)separating starburst galaxies and AGN on all three diagnostic plots,or(2)fell within the domain of starburst galaxies in one diagnostic diagram and AGN in another.The sample was constrained by a redshift upper limit of z=0.035and declinations south of+24◦.From the optical spectroscopic surveys of Veilleux&Osterbrock(1987),van den Brock et al.(1991),Ashby et al.(1995)and Veilleux et al.(1995),we identiﬁed34galaxies in which the nuclear spectra have ambiguous optical line diagnostics satisfying the above criteria.NIR spectroscopy has been obtained for12of these galaxies and for a furtherﬁve galaxies that were observed for comparison,comprising one starburst and four AGNs.2.3.Modelling the Optical DiagnosticsWhile the segregation of starbursts and Seyferts within the optical line ratio diagnostic diagrams was discovered observationally,photoionisation models(Evans&Dopita1985),using hot stars and a powerlaw input spectrum to model starbursts and AGN respectively, conﬁrm the observationalﬁnding.However,in the time since these models were calculated, knowledge of atomic structure and metal abundances has improved.Therefore,we use the latest photoionisation code MAPPINGS II(see Sutherland&Dopita1993)to examine the dependence of the optical diagnostic line ratios on metallicity(Model A),stellar temperature and hardness of the powerlaw spectrum(Model B)and hydrogen density (Model C),according to the parameters outlined in Table1.MAPPINGS II has also been newly expanded so that it now includes modelling of NIR emission lines(see Section4).The stellar atmosphere models used in MAPPINGS II are from Hummer&Mihalas (1970)and the AGN models are created from a powerlaw ionising spectrum of the form fν∼να.The ionisation parameter,U,is deﬁned by U=(n Hǫ2l C)1/3,where n H is the total number density of hydrogen atoms and ions,l C is the number of hydrogen-ionising photons emitted per unit time from the central source,andǫis the volumeﬁlling factor(unity in our models).This parameter was varied from log U=−3to0for starburst models and log U=−4to−1for AGN models.Spherical geometry and isochoric conditions were assumed throughout.The calculations were terminated when the fraction of neutral hydrogen reached 95%for both the starburst and AGN powerlaw models,thus permitting a uniform outer boundary condition.While a starburst model could be terminated at a higher neutral fraction,the partially ionised AGN models have numerical diﬃculties in achieving complete neutrality,since the thermal balance becomes due solely to X-ray ionisation and electron cascade heating.At a95%neutral fraction,the overwhelming majority of optical and NIR emission from the models has occurred,so no signiﬁcant error is introduced.The AGN models may not be accurate for far infraredﬁne structure lines,but these are not considered in this paper.2.3.1.Model LimitationsA simplifying geometric assumption has been built into the models,namely that photoionisation occurs within a single homogeneous sphere.As a result,the partially ionised region produced by the models exists as a very thin,outer shell.In reality,the interstellar medium(ISM)is not homogeneous and a thin shell of partially ionised gas will form around each density clump found within an ionisation region.This creates a larger eﬀective emission volume for the[S II]and[O I]lines,that are formed predominantly in partially ionised regions.As the ionisation parameter(U)is increased within the models, corresponding to stronger[O III]emission,the partially ionised region decreases moresharply than expected.This is because a non-homogeneous medium will contain clumps of varying density,so that even as the ionisation parameter increases the total volume of partially ionised gas can remain approximately constant.This eﬀect causes the modelsto slightly underestimate emission from the partially ionised region,particularly[S II] emission.It is most pronounced in starburst models,where the partially ionised region is necessarily small due to the relatively soft radiation from stellar sources,compared with the harder ionising spectrum of a powerlaw model.Modelling of[O I]is further limited by the complexity of the[O I]emission process. The collision strengths for transitions of the neutral[O I]species are a strong functionof temperature.Further,collisions by protons and neutral hydrogen atoms can make important contributions to[O I]emission,especially in partially ionised zones that occur in AGN.However,the rates for proton and neutral hydrogen collisions are not very well known,making accurate predictions diﬃcult for[O I]in AGN(see Williams&Livio1995 for reviews and discussions on atomic data for emission lines).In starburst models,where the[O I]is conﬁned to a very thin layer in the outer radii,the[O I],and to a lesser extent [S II],emission is sensitive to the chosen boundary conditions.On the other hand,the[O III]/Hβand[N II]/Hαratios are comparatively more reliable,from the atomic data and modelling point of view.MAPPINGS II uses a six level atomic model for O++and N+.Electron collisions alone are the primary excitation mode for the forbidden lines and the temperature dependence of the collision strengths of these ions is relatively weak and well determined.The caveat,however,is that the[N II]/Hαratio is diﬃcult to measure at low spectral resolution due to the blending of the[N II]and Hαemission lines.2.3.2.Model AnalysisThe results for the three diﬀerent models,deﬁned in Table1,are shown in Figures1–3,with the starburst models denoted by‘s’and the powerlaw models denoted by‘p’.The observational data are taken from Veilleux&Osterbrock(1987,and papers therein)and we also include our sample of composite galaxies.Figure1shows the eﬀect of variation in metallicity.Low and intermediate metallicities are based on the abundances of the SMC and the LMC respectively(Russell&Dopita 1990).The abundances for solar metallicity are taken from Anders&Grevesse(1989)and in a further variation,the solar abundances were depleted using the IUE data of Shull (1993).The interesting result evident from these models occurs within the diagnostic diagram of[O III]/Hβversus[N II]/Hα(Figure1a).It can be seen that when the metallicity is below solar the powerlaw models(i.e.p1and p2)fall into the domain of starbursts.This suggests that mis-classiﬁcation of narrow emission line galaxies using the diagnostic diagrams could occur for low metallicity AGN.Figure2shows the eﬀect of stellar temperature on the starburst models and the hardness of the ionising spectrum on the AGN models.The diagrams suggest that the starburst galaxies are best modelled by high temperatures(40000K–45000K),while the AGNs require a powerlaw indexα=−1.5to−2.0.Figure3shows the eﬀect of varying the hydrogen density.The powerlaw model does not converge for n H=102cm−3,and therefore is not included in theﬁgure.As expected, starburst galaxies are bestﬁtted by models with n H=103to104cm−3,typical of the H II regions in the Orion Nebula.A change in hydrogen density has little eﬀect on the powerlaw models.Overall,the starburst and powerlaw models follow the segregation of the starburst and AGN data,with the composite galaxies consistently falling between the two model types.3.Observations and ReductionsEchelle spectroscopy was obtained over the period1994–1996using the InfraRed Imaging Spectrograph(IRIS)at the f/36Cassegrain focus of the3.9m Anglo-Australian Telescope(AAT).Two echelles were used,the IJ echelle(0.9−1.5µm)and the HK echelle (1.46−2.5µm),each with spectral resolutionλ/∆λ≈400.A non-destructive readout method,with read noise around40e−rms,was used,whereby the array is sampled regularly during an integration.This method has the advantage that data acquired after the saturation level is reached are removed.The log of the observations is presented in Table2. Flux standards of spectral classes A and G were observed at zenith distances similar to the target galaxies.Wavelength calibration was performed using comparison lamps of helium, argon and xenon.The galaxy spectra were reduced with the STARLINK program FIGARO using subroutines written speciﬁcally for IRIS.The observations were performed in pairs,with the galaxy positioned alternately at each end of the13′′-long slit.Subtraction of these image pairs provides a goodﬁrst order sky subtraction and any residual sky is removed when pairs of extracted spectra are averaged together.Aﬂat-ﬁeld image was formed by observing the dome windscreen illuminated by a tungsten lamp;a second exposure with the lamp oﬀwasthen subtracted to remove thermal radiation from the telescope and surroundings.The ﬂat-ﬁeld image was used to remove wavelength dependent pixel-to-pixel variations across the array.Individual bad pixels were removed from the images by linear interpolation. Straightening of the echelle orders was achieved using standard star frames,while the comparison lamp spectra were used to correct for non-vertical positioning of the slit.Before using the standard stars toﬂux calibrate the galaxy spectra and correct for the atmosphere, hydrogen absorption lines intrinsic to the standard were removed.Each standard star spectrum was thenﬁtted to a blackbody model with a temperature appropriate for its spectral class.It is important that approximately the same spatial region be examined for all the galaxies,to provide consistency in the results.We chose,therefore,to produce a nuclear spectrum by extracting pixels over a window corresponding to a linear scale of≈1kpc at the redshift of each galaxy.For those galaxies near the redshift upper limit of the sample (z=0.035),a linear scale of1kpc corresponds to a2pixel extraction box.IRIS has a pixel scale of0.79′′/pixel and therefore the extraction box matches the average seeing of1.5′′, obtained during the observations.The individual orders were combined into one spectrum covering the entire IJ and HK bands.This involved bringing all the orders to the same pixel scale(≈30˚A/pixel).The spectra were shifted to rest wavelength using published redshifts.Spectra of the sample of composite galaxies are shown in Figure4and the comparison spectra of starbursts and AGN are shown in Figure5.The emission lineﬂuxes were measured using the IRAF routine SPLOT,whichﬁts a gaussian to the line proﬁles.The relative strengths of the emission lines are given in Tables3and4with measurement errors estimated to be≈10%.The composite galaxies with both IJ and HK data were used to form a co-added spectrum,which is shown in Figure6.The individual galaxy spectra,all of similar slopes, wereﬁrst normalised within the region1.5–1.6µm and then averaged together.Theco-added spectrum clearly brings out the emission features detected within the sample.4.Results and Discussion4.1.Extinction CorrectionsExtinction corrections in the optical were calculated following Veilleux&Osterbrock (1987),using the Whitford reddening curve parametrized by Miller&Mathews(1972)(seeTable5).The intrinsic Hα/Hβratio is usually set at2.85for starburst galaxies(assuming Case B recombination)and3.10for AGN,reﬂecting enhanced Hαemission due to collisional excitation.For our sample,where we are unsure of the classiﬁcation,we found that neither value resolved the ambiguous nature of the line ratios and chose to adopt Hα/Hβ=2.85 for the composite galaxies.Extinction in the NIR can be calculated using the NIR hydrogen recombination lines of Paβand Brγ.However,with IRIS these two lines are recorded using diﬀerent echelles, Paβwithin IJ and Brγwithin HK.A valid comparison requires photometric conditions during observations in both wavebands.This occurred for only three galaxies,Mrk52, ESO602-G025and NGC7130and the values derived for the extinction from the NIR lines, E(B–V)NIR,are0.0,1.17and1.23,respectively.These values for E(B–V)NIR are consistent with those based on the optical line ratios for the three galaxies(see Table5)implying that the dust is found within a homogeneous foreground screen,as opposed to it being mixed with the line-emitting gas(Puxley&Brand1994;Calzetti et al.1996).4.2.Emission Features4.2.1.[Fe II]EmissionIn the same manner as the optical forbidden lines,[Fe II]λλ1.25,1.64is produced within regions that are partially ionised,suggesting that[Fe II],like[O I]for example,might be a useful indicator of AGN activity.However,[Fe II]can also be produced via shock excitation.In fact,strong[Fe II]emission is found within starburst galaxies,most likely a result of shock excitation by SNRs(Mouri et al.1993;Forbes&Ward1993;Vanzi&Rieke 1997),with a possible contribution from partially ionised regions created within cooling zones formed behind the shock front in SNRs(Oliva et al.1989).SNRs may also contribute to the[Fe II]emission in AGN and,in addition,shock excitation can occur through the interaction of possible jets and outﬂows from the AGN with the surrounding medium.There is some indication that[Fe II]emission may be correlated with radio emission in starbursts and AGNs(Forbes&Ward1993).Within starbursts the correlation is readily attributed to shock excitation from SNRs.Since both starbursts and AGNs follow the same correlation it lends support to the idea that SNRs may be producing shock excitation within AGN.However,in at least some AGNs the shock excitation is more likely associated with jet interactions,as conﬁrmed in NGC1068(Blietz et al.1994)where the[Fe II]emission appears to trace the radio jet.To further complicate matters,Fe is readily depleted onto dust grains and so thedestruction of dust grains by shocks can lead to an enhancement of[Fe II]as it is released into the gas phase(Greenhouse et al.1991).4.2.2.[Fe II]/Paβvs[O I]/HαThe optical line ratio,[O I]/Hα,is established as a reliable diagnostic for segregating AGN from starbursts(Veilleux&Osterbrock1987),while the NIR line ratio, [Fe II](1.25µm)/Paβ,has also shown promise as a good diagnostic(Simpson et al.1996; Alonso-Herrero et al.1997)and has a similar advantage of utilising lines close in wavelength, so reddening eﬀects are minimised.In Figure7we present a plot of[Fe II]/Paβversus[O I]/Hαfor our sample of composite galaxies,along with data from the literature(Mouri et al.1990,1993;Simpson et al.1996) and photoionisation model predictions from MAPPINGS II.All the galaxies lie within the region spanned by the Orion Nebula and SNRs,which may indicate a progression from pure photoionisation to pure shock excitation(Alonso-Herrero et al.1997).However,starburst galaxies have been found to lie along the mixing curve combining H II regions and SNRs, whereas Seyfert galaxies do not(Simpson et al.1996).Our sample galaxies occupy the region between starbursts and Seyferts,further supporting their composite nature.The same photoionisation models used in Section2.3are included in Figure7,showing the eﬀect of varying the metallicity,stellar temperature,powerlaw index and hydrogen density.Our results are consistent with the CLOUDY models of Alonso-Herrero et al. (1997),including theﬁnding that the AGN data are bestﬁtted by powerlaw models with low metallicity.However,Alonso-Herrero et al.(1997)ﬁnd little diﬀerence between models with and without grains,using grain properties derived from the Orion Nebula.Our dust depletion model is based on depletion factors in the local ISM(Shull1993),applied to a solar abundance(Anders&Grevesse1989).Weﬁnd this model to be too dusty to explain the observed[Fe II]/Paβratio,a reﬂection of the fact that Fe is very heavily depleted in the local ISM.As outlined in Section4.2.1,there is much evidence that shock excitation isthe dominant mechanism for producing[Fe II]emission.Thus,it is surprising thatthe photoionisation modelsﬁt the data so well.Shock models of these ratios using MAPPINGS II are under development(Sutherland1998).Two galaxies,Mrk52and MCG-02-33-098,have considerably lower[Fe II]/Paβratios than the rest of the sample.In fact,Figure7provides strong evidence that Mrk52is dominated by star formation,aﬁnding that is further supported by strong CO absorptionand Brγemission(see Section4.3.2).MCG-02-33-098is an interacting system showing severely disturbed morphology.Of the two nuclei detected,it is the western nucleus which has a composite optical spectrum while the eastern nucleus has optical line ratios consistent with a starburst galaxy.The merger itself may be responsible for the unusual line ratios, for example by increasing shock excitation by cloud-cloud collisions.However,when we examine the optical spectrum for this galaxy(Veilleux et al.1995)it appears quite noisy and the[O I]/Hαratio may have been overestimated.If so,MCG-02-33-098may move closer to the other starburst galaxies in Figure7.4.2.3.[Fe II]/Brγvs H2/BrγThe usefulness of optical emission line diagnostic diagrams has prompted the search for similar diagrams using NIR lines.The lower dust extinction in the NIR is an advantage in probing close to the nucleus.Moorwood and Oliva(1988)proposed[Fe II](1.644µm)/Brγversus H2(2.122µm)/Brγas a possible NIR diagnostic tool for distinguishing starburst and AGN emission,not least because these are the strongest and,therefore,the most easily measured lines.Their data suggest a segregation of starburst galaxies towards the region of low[Fe II]/Brγand H2/Brγwhich can readily be explained by the enhanced strength of Brγin starburst galaxies compared with AGN(cf.Section2.1).For completeness we present the NIR line ratio diagram of[Fe II]/Brγversus H2/Brγin Figure8.Our sample of galaxies is plotted together with the original measurements of Moorwood&Oliva(1988),data compiled by Forbes&Ward(1993)and blue dwarf galaxies from Vanzi&Rieke(1997).The composite galaxies,like the starbursts and Seyfert2s,span the range from blue dwarfs to Seyfert1s.Unfortunately,there is no strong separation of starbursts and Seyfert2s,limiting the usefulness of this diagram as a diagnostic.4.2.4.[S III]/Paβvs[S II]/HαEmission line ratios involving sulfur have been found useful for distinguishing between photoionisation and shock excitation(Diaz,Pagel&Terlevich1985;Diaz,Pagel&Wilson 1985).If shock excitation is occurring,it is expected that the[S III]line should be weaker than[S II]since S++,in shocked gas,cools predominantly via UV line emission, corresponding to higher temperatures,rather than by emission in the NIR(Dopita1977).In Figure9we present a comparison of the sulfur line ratios[S III]/Paβand[S II]/Hα. There is a slight tendency for the[S III]/Paβratios of the composite galaxies to favour astarburst origin as they all lie in the region of log([S III]/Paβ)<1which is favoured by starbursts.On the other hand,the[S II]/Hαline ratio is used in the diagnostic diagrams of Veilleux&Osterbrock(1987)and therefore the composite galaxies have been chosen speciﬁcally so that they lie intermediate between the starbursts and AGN.SNRs are also included in Figure9and fall in the region of low[S III]emission,as predicted for shock excitation.The usual photoionisation models are also shown in Figure9.The powerlaw models are a goodﬁt to the AGN data.However,the starburst models are a poorﬁt,with the [S III]/Paβratios being too high overall toﬁt the starburst data.As discussed above, shocks can weaken the[S III]/Paβratio,hinting that shock excitation may be an important emission mechanism within the starburst galaxies.4.3.Absorption Features4.3.1.CO IndicesProduction of CO occurs in the outer envelopes of late-type stars,so the CO absorption bandhead longward of2.3µm appears prominently in the spectra of red giant and supergiant stars.Evolutionary modelling(Doyon et al.1994)of starburst galaxies predicts a sharp increase in CO band strength after approximately107yr,as red supergiants appear in the stellar population.In principle,therefore,the observed CO absorption is a means for constraining the age of a starburst.It has only been in recent years that NIR spectroscopic detections of CO absorption in galaxies have been made and there is no clear consensus regarding the best method for measuring the strength of the absorption.The earliest CO detections,from low resolution data,show a single absorption feature longwards of2.3µm.The original spectroscopic CO index(CO sp)is deﬁned asCO sp=−2.5log10( R2.36 )(1) where R2.36 is the average CO depth of the rectiﬁed spectrum between2.31and2.40µm (Doyon et al.1994).The rectiﬁed spectrum is obtained byﬁtting a powerlaw(fλ∝λβ)to featureless regions of the continuum between2.00and2.29µm and extrapolating to longer wavelengths.CO sp was deﬁned in this manner to be compatible with the photometric CO index(CO ph)measured using narrowbandﬁlters having eﬀective wavelengths(λe)and FWHM(∆λ)ofλe=2.20,∆λ=0.11µm(continuum)andλe=2.36,∆λ=0.08µm(CO absorption),respectively(Frogel et al.1978).A conversion between CO sp and CO ph is given by Doyon et al.(1994)based on the NIR spectroscopic stellar atlas of Kleinmann&Hall。