XMM-Newton and Broad Iron Lines

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Comparing Suzaku and XMM-Newton Observations of the Soft X-ray Background Evidence for Sola

a r X i v :0712.3538v 1 [a s t r o -p h ] 20 D e c 2007Draft version February 2,2008Preprint typeset using L A T E X style emulateapj v.10/09/06COMPARING SUZAKU AND XMM-NEWTON OBSERVATIONS OF THE SOFT X-RAY BACKGROUND:EVIDENCE FOR SOLAR WIND CHARGE EXCHANGE EMISSIONDavid B.Henley and Robin L.SheltonDepartment of Physics and Astronomy,University of Georgia,Athens,GA 30602Draft version February 2,2008ABSTRACTWe present an analysis of a pair of Suzaku spectra of the soft X-ray background (SXRB),obtained from pointings on and oﬀa nearby shadowing ﬁlament in the southern Galactic hemisphere.Because of the diﬀerent Galactic column densities in the two pointing directions,the observed emission from the Galactic halo has a diﬀerent shape in the two spectra.We make use of this diﬀerence when modeling the spectra to separate the absorbed halo emission from the unabsorbed foreground emission from the Local Bubble (LB).The temperatures and emission measures we obtain are signiﬁcantly diﬀerent from those determined from an earlier analysis of XMM-Newton spectra from the same pointing directions.We attribute this diﬀerence to the presence of previously unrecognized solar wind charge exchange (SWCX)contamination in the XMM-Newton spectra,possibly due to a localized enhancement in the solar wind moving across the line of sight.Contemporaneous solar wind data from ACE show nothing unusual during the course of the XMM-Newton observations.Our results therefore suggest that simply examining contemporaneous solar wind data might be inadequate for determining if a spectrum of the SXRB is contaminated by SWCX emission.If our Suzaku spectra are not badly contaminatedby SWCX emission,our best-ﬁtting LB model gives a temperature of log(T LB /K)=5.98+0.03−0.04and a pressure of p LB /k =13,100–16,100cm −3K.These values are lower than those obtained from other recent observations of the LB,suggesting the LB may not be isothermal and may not be in pressure equilibrium.Our halo modeling,meanwhile,suggests that neon may be enhanced relative to oxygen and iron,possibly because oxygen and iron are partly in dust.Subject headings:Galaxy:halo—Sun:solar wind—X-rays:diﬀuse background—X-rays:ISM1.INTRODUCTIONThe diﬀuse soft X-ray background (SXRB),which is observed in all directions in the ∼0.1–2keV band,is composed of emission from several diﬀerent components.For many years,the observed 1/4-keV emission was believed to originate from the Local Bubble (LB),a cavity in the local interstellar medium (ISM)of ∼100pc radius ﬁlled with ∼106K gas (Sanders et al.1977;Cox &Reynolds 1987;McCammon &Sanders 1990;Snowden et al.1990).However,the discovery of shadows in the 1/4-keV background with ROSAT showed that ∼50%of the 1/4keV emission originates from beyond the LB (Burrows &Mendenhall 1991;Snowden et al.1991).This more distant emission origi-nates from the Galactic halo,which contains hot gas with temperatures log(T halo /K)∼6.0–6.5(Snowden et al.1998;Kuntz &Snowden 2000;Smith et al.2007;Galeazzi et al.2007;Henley,Shelton,&Kuntz 2007).As the halo gas is hotter than the LB gas,it also emits at higher energies,up to ∼1keV.Above ∼1keV the X-ray background is extragalactic in origin,and is due to un-resolved active galactic nuclei (AGN;Mushotzky et al.2000).X-ray spectroscopy of the SXRB can,in principle,de-termine the thermal properties,ionization state,and chemical abundances of the hot gas in the Galaxy.These properties give clues to the origin of the hot gas,which is currently uncertain.However,to de-termine the physical properties of the hot Galactic gas,one must ﬁrst decompose the SXRB into its con-Electronic address:dbh@stituents.This is achieved using a technique called “shadowing”,which makes use of the above-mentioned shadows cast in the SXRB by cool clouds of gas be-tween the Earth and the halo.Low-spectral-resolution ROSAT observations of the SXRB were decomposed into their foreground and background components by modeling the intensity variation due to the varying absorption column density on and around shadow-ing clouds (Burrows &Mendenhall 1991;Snowden et al.1991,2000;Snowden,McCammon,&Verter 1993;Kuntz,Snowden,&Verter 1997).The same tech-nique was used to decompose ROSAT All-Sky Survey data over large areas of the sky (Snowden et al.1998;Kuntz &Snowden 2000).With higher resolution spectra,such as those from the CCD cameras onboard XMM-Newton or Suzaku ,it is possible to decompose the SXRB into its foreground and background components spectroscopically.This is achieved using one spectrum toward a shadowing cloud,and one toward a pointing to the side of the cloud.The spectral shape of the absorbed background component (and hence of the overall spectrum)will diﬀer between the two directions,because of the diﬀerent absorbing col-umn densities.Therefore,by ﬁtting a suitable multicom-ponent model simultaneously to the two spectra,one can separate out the foreground and background emis-sion components.Such a model will typically consist of an unabsorbed single-temperature (1T )thermal plasma model for the LB,an absorbed thermal plasma model for the Galactic halo,and an absorbed power-law for the ex-tragalactic background.The Galactic halo model could be a 1T model,a two-temperature (2T )model,or a dif-2HENLEY AND SHELTONferential emission measure (DEM)model (Galeazzi et al.2007;Henley et al.2007;S.J.Lei et al.,in preparation).Recent work has shown that there is an additional complication,as X-ray emission can originate within the solar system,via solar wind charge exchange (SWCX;Cox 1998;Cravens 2000).In this process,highly ion-ized species in the solar wind interact with neutral atoms within the solar system.An electron transfers from a neutral atom into an excited energy level of a solar wind ion,which then decays radiatively,emitting an X-ray photon.The neutral atoms may be in the outer reaches of the Earth’s atmosphere (giving rise to geo-coronal emission),or they may be in interstellar ma-terial ﬂowing through the solar system (giving rise to heliospheric emission).It has been estimated that the heliospheric emission may contribute up to ∼50%of the observed soft X-ray ﬂux (Cravens 2000).The geocoro-nal emission is typically an order of magnitude fainter,but during solar wind enhancements it can be of similar brightness to the heliospheric emission (Wargelin et al.2004).SWCX line emission has been observed with Chandra ,XMM-Newton ,and Suzaku (Wargelin et al.2004;Snowden,Collier,&Kuntz 2004;Fujimoto et al.2007).As the SWCX emission is time varying,it cannot easily be modeled out of a spectrum of the SXRB.If SWCX contamination is not taken into account,analyses of SXRB spectra will yield incorrect results for the LB and halo gas.This paper contains a demonstra-tion of this fact.Henley et al.(2007)analyzed a pair of XMM-Newton spectra of the SXRB using the pre-viously described shadowing technique.One spectrum was from a direction toward a nearby shadowing ﬁlament in the southern Galactic hemisphere (d =230±30pc;Penprase et al.1998),while the other was from a direc-tion ∼2◦away.The ﬁlament and the shadow it casts in the 1/4-keV background are shown in Figure 1.Contem-poraneous solar wind data from the Advanced Composi-tion Explorer (ACE )showed that the solar wind was steady during the XMM-Newton observations,without any ﬂares or spikes.The proton ﬂux was slightly lower than average,and the oxygen ion ratios were fairly typ-ical.These observations led Henley et al.(2007)to con-clude that their spectra were unlikely to be severely con-taminated by SWCX ing a 2T halo model,they obtained a LB temperature of log(T LB /K)=6.06and halo temperatures of log(T halo /K)=5.93and 6.43.The LB temperature and the hotter halo temperature are in good agreement with other recent measurements of the SXRB using XMM-Newton and Suzaku (Galeazzi et al.2007;Smith et al.2007),and with analysis of the ROSAT All-Sky Survey (Kuntz &Snowden 2000).We have obtained spectra of the SXRB from the same directions as Henley et al.’s (2007)XMM-Newton spectra with the X-ray Imaging Spectrometer (XIS;Koyama et al.2007)onboard the Suzaku X-ray obser-vatory (Mitsuda et al.2007).The XIS is an excellent tool for studying the SXRB,due to its low non-X-ray background and good spectral resolution.Our Suzaku pointing directions are shown in Figure 1.We analyze our Suzaku spectra using the same shadowing technique used by Henley et al.(2007).We ﬁnd that there is poor agreement between the results of our Suzaku analysis and the results of the XMM-Newton analysis in Henley et al.(2007).We attribute this discrepancy to previously un-Fig. 1.—The shadowing ﬁlament used for our observations,shown in Galactic coordinates.Grayscale :ROSAT All-Sky Sur-vey R1+R2intensity (Snowden et al.1997).Contours :IRAS 100-micron intensity (Schlegel et al.1998).Yellow squares :Our Suzaku pointing directions.recognized SWCX contamination in the XMM-Newton spectra,which means that SWCX contamination can oc-cur at times when the solar wind ﬂux measured by ACE is low and does not show ﬂares.This paper is organized as follows.The Suzaku obser-vations and data reduction are described in §2.The anal-ysis of the spectra using multicomponent spectral models is described in §3.The discrepancy between the Suzaku results and the XMM-Newton results is discussed in §4.This discrepancy is due to the presence of an additional emission component in the XMM-Newton spectra,which we also describe in §4.In §5we measure the total intensi-ties of the oxygen lines in our Suzaku and XMM-Newton spectra.We concentrate on these lines because they are the brightest in our spectra,and are a major component of the 3/4-keV SXRB (McCammon et al.2002).In §6we present a simple model for estimating the intensity of the oxygen lines due to SWCX,which we compare with our observations.We discuss our results in §7,and con-clude with a summary in §8.Throughout this paper we quote 1σerrors.2.OBSERVATIONS AND DATA REDUCTIONBoth of our observations were carried out in early 2006March.The details of the observations are shown in Ta-ble 1.In the following,we just analyze data from the back-illuminated XIS1chip,as it is more sensitive at lower energies than the three front-illuminated chips.Our data were initially processed at NASA Goddard Space Flight Center (GSFC)using processing version 1.2.2.3.We have carried out further processing and ﬁl-tering,using HEAsoft 1v6.1.2and CIAO 2v3.4.We ﬁrst combined the data taken in the 3×3and 5×5observa-tion modes.We then selected events with grades 0,2,3,4,and 6,and cleaned the data using the standard data selection criteria given in the Suzaku Data Reduc-tion Guide 3.We excluded the times that Suzaku passed through the South Atlantic Anomaly (SAA),and also times up to 436s after passage through the SAA.We also excluded times when Suzaku ’s line of sight was el-evated less than 10◦above the Earth’s limb and/or was1/lheasoft 2/ciao3/docs/suzaku/analysis/abc/abc.htmlOBSERVING CHARGE EXCHANGE WITH SUZAKU AND XMM-NEWTON3TABLE1Details of Our Suzaku ObservationsObservation l b Start time End time Usable exposureID(deg)(deg)(UT)(UT)(ks)Oﬀﬁlament501001010278.71−47.072006-03-0116:56:012006-03-0222:29:1455.6Onﬁlament501002010278.65−45.302006-03-0320:52:002006-03-0608:01:1969.0less than20◦from the bright-Earth terminator.Finally,we excluded times when the cut-oﬀrigidity(COR)wasless than8GV.This is a stricter criterion than that inthe Data Reduction Guide,which recommends exclud-ing times with COR<6GV.However,the higher CORthreshold helps reduce the particle background,and forobservations of the SXRB one desires as low a particlebackground as possible.The COR threshold that weuse has been used for other Suzaku observations of theSXRB(Fujimoto et al.2007;Smith et al.2007).Finally,we binned the2.5–8.5keV data into256-s time bins,andused the CIAO script analyzepublic4HENLEY AND SHELTON34:00.03:33:00.032:00.031:00.015:00.020:00.025:00.0-63:30:00.035:00.040:00.0Right ascensionD e c l i n a t i o nOn filament22:00.021:00.03:20:00.019:00.015:00.0-62:20:00.025:00.030:00.035:00.0Right ascensionD e c l i n a t i o nOff filamentFig.2.—Cleaned and smoothed Suzaku XIS1images in the 0.3–5keV band for our on-ﬁlament (left )and oﬀ-ﬁlament (right )observations.The data have been binned up by a factor of 4in the detector’s x and y directions,and then smoothed with a Gaussian whose standard deviation is equal to 1.5times the binned pixel size.The particle background has not been subtracted from the data.The red circles outline the regions that were excluded from the analysis (see text for details).in our ﬁt (Snowden et al.1997).The R1and R2count-rates help constrain the model al lower energies (below ∼0.3keV),while the higher channels overlap in energy with our Suzaku spectra.We extracted the ROSAT spec-tra from 0.5◦radius circles centered on our two Suzaku pointing directions using the HEASARC X-ray Back-ground Tool 8v2.3.The spectral analysis was carried out using XSPEC 9v11.3.2(Arnaud 1996).For the thermal plasma com-ponents,we used the Astrophysical Plasma Emission Code (APEC)v1.3.1(Smith et al.2001)for the Suzaku data and the ROSAT R4–7bands,and the Raymond-Smith code (Raymond &Smith 1977and updates)for the ROSAT R1–3bands.For a given model component (i.e.,the LB or one of the two halo components),the tem-perature and normalization of the ROSAT Raymond-Smith model are tied to those of the corresponding Suzaku APEC model.We chose to use the Raymond-Smith code for the lower-energy ROSAT channels be-cause APEC’s spectral calculations below 0.25keV are inaccurate,due to a lack of data on transitions from L-shell ions of Ne,Mg,Al,Si,S,Ar,and Ca 10.As the upper-limit of the ROSAT R1and R2bands is 0.284keV,and the R3band also includes such low-energy photons (Snowden et al.1997),APEC is not ideal for ﬁtting to these energy bands.For the absorption,we used the XSPEC phabs model,which uses cross-sections from Ba l uci´n ska-Church &McCammon (1992),with an updated He cross-section from Yan,Sadeghpour,&Dalgarno (1998).Following Henley et al.(2007),we used the interstellar chemical abundance table from Wilms,Allen,&McCray (2000).For many astrophysically abundant elements,these abundances are lower than those in the widely used solar abundance table of Anders &Grevesse (1989).However,recently several elements’solar photospheric abundances have been revised downwards (Asplund et al.2005),and are in good agreement with the Wilms et al.(2000)8/cgi-bin/Tools/xraybg/xraybg.pl9/docs/xanadu/xspec/xspec1110/atomdb/issuesOBSERVING CHARGE EXCHANGE WITH SUZAKU AND XMM-NEWTON5-4-2 0 2 40.3135channel energy (keV)(d a t a -m o d e l )/σ0.001 0.010.11c o u n t s s -1k e V-1On filamentLocal Bubble Halo (cold)Halo (hot)ExtragalacticInstrumental lines TotalO V I I O V I I IN e I XM g X I ?A l KS i KA u M -4-2 0 2 40.3135channel energy (keV)(d a t a -m o d e l )/σ0.0010.010.11c o u n t s s -1k e V-1Off filamentLocal Bubble Halo (cold)Halo (hot)ExtragalacticInstrumental lines TotalN V I I ?O V I I O V I I IN e I XA l KS i KA uMFig.3.—Our observed on-ﬁlament (left )and oﬀ-ﬁlament (right )Suzaku spectra,with the best-ﬁtting model obtained by ﬁtting jointly to the Suzaku and ROSAT data (Model 1in Table 7).The gap in the Si K instrumental line is where channels 500–504have been removed from the data (see §2).10100100010-6 c o u n t s s -1 a r c m i n-2-4-2 0 2 4R1R2R3R4R5R6R7ROSAT band(d a t a - m o d e l ) / σFig. 4.—The on-ﬁlament (dashed )and oﬀ-ﬁlament (solid )ROSAT All-Sky Survey spectra,compared with Model 1from Ta-ble 7.For clarity the individual model components have not been plotted.CALDB.This model component attenuated the emis-sion from the LB,halo,and extragalactic background for the Suzaku spectra only.We adjusted the model oxygen abundance to give C /O =6(Koyama et al.2007),and set the abundances of all other elements to zero.The results of this model are given as Model 2in Table 7.Figures 5and 6show this model com-pared with the Suzaku and ROSAT spectra,respec-tively.One can see from Figure 6that the ﬁt to the ROSAT data is greatly improved.The model implies a column density of carbon atoms,N C =(0.28±0.04)×1018cm −2,in addition to the amount of contamina-tion given by the CALDB contamination model,which is N C =3.1×1018cm −2at the center of the XIS1chip(from the CALDB ﬁle aecontami xrtxis08.html6HENLEY AND SHELTON-4-2 0 2 40.3135channel energy (keV)(d a t a -m o d e l )/σ0.0010.010.1c o u n t s s -1k e V-1On filamentLocal Bubble Halo (cold)Halo (hot)ExtragalacticInstrumental lines Total-4-2 0 2 40.3135channel energy (keV)(d a t a -m o d e l )/σ0.0010.010.1c o u n t s s -1k e V-1Off filamentLocal Bubble Halo (cold)Halo (hot)ExtragalacticInstrumental linesTotalFig. 5.—As Figure 3,but with a vphabs component included for the Suzaku spectra to model XIS contamination above that included in the CALDB (Model 2in Table 7;see §3.2for details).10 100100010-6 c o u n t s s -1 a r c m i n-2-4-2 0 2 4R1R2R3R4R5R6R7ROSAT band(d a t a - m o d e l ) / σFig.6.—As Figure 5,but for Model 2from Table 7.In Model 5we ﬁt exactly the same model to the Suzaku data alone.Without the ROSAT data,we cannot con-strain the LB temperature T LB .We therefore ﬁx it at the value determined in Model 2:T LB =105.98K.We also ﬁx N C for the vphabs contamination component at the Model 2value:N C =0.28×1018cm −2.The best-ﬁtting model parameters are in very good agree-ment with those obtained by ﬁtting jointly to the Suzaku and ROSAT spectra (compare Models 2and 5).We can only get consistent results between the Suzaku -ROSAT joint ﬁt and the ﬁt to just the Suzaku data by using a two-temperature model for the halo.However,note that the Model 5LB emission measure is consistent (within its errorbar)with zero.This is not to say that our data imply that there is no LB at all:as noted above,we need a LB component and two halo components to get a good joint ﬁt to the Suzaku and ROSAT data.Instead,the Model 5results imply that our Suzaku data are consis-tent with the LB not producing signiﬁcant emission in the Suzaku band (i.e.,E 0.3keV).Also shown in Table 7are the results of ﬁtting our model (without any LB component)to the on-and oﬀ-ﬁlament Suzaku spectra individually (Models 6and 7,respectively).As has already been noted,the normaliza-tion of the extragalactic background diﬀers signiﬁcantly between the two spectra.However,the plasma model pa-rameters for the individual spectra are in good agreement with each other.These results seem to justify allowing the extragalactic normalization to diﬀer between the two spectra while keeping all other model parameters equal for the two spectra.3.3.Chemical Abundances in the HaloWe investigated the chemical abundances of the X-ray-emissive halo gas by repeating the above-described mod-eling,but allowing the abundances of certain elements in the halo components to vary.In particular we wished to investigate whether or not varying the neon and magne-sium halo abundances improved the ﬁts to the Ne ix and Mg xi features noted above.We also allowed the abun-dance of iron to vary,as iron is an important contributor of halo line emission to the Suzaku band.For this investigation we just ﬁt to the Suzaku data,ﬁxing the LB temperature at log(T LB /K)=5.98,and ﬁxing the carbon column density of the vphabs contam-ination component at N C =0.28×1018cm −2.As we could not accurately determine the level of the contin-uum due to hydrogen,we could not measure absolute abundances.Instead,we estimated the abundances rel-ative to oxygen by holding the oxygen abundance at its Wilms et al.(2000)value,and allowing the abundances of neon,magnesium,and iron to vary.Both halo com-ponents were constrained to have the same abundances.The best-ﬁtting temperatures and emission measures of the various model components are presented as Model 8in Table 7,and the abundances are presented in Table 2.The best-ﬁtting model parameters are not signiﬁcantly aﬀected by allowing certain elements’abun-dances to vary (compare Model 8with Model 5).Iron does not seem to be enhanced or depleted relative to oxy-gen in the halo.Neon and magnesium both appear to be enhanced in the halo relative to oxygen,which is what one would expect from Figures 3and 5,as the modelsOBSERVING CHARGE EXCHANGE WITH SUZAKU AND XMM-NEWTON7 TABLE2Halo AbundancesElement AbundanceO1(ﬁxed)Ne a1.8±0.4Mg a4.6+3.5−2.8Fe1.2+0.4−0.5Note.—Abun-dances are relative tothe Wilms et al.(2000)interstellar abundances:Ne/O=0.178,Mg/O=0.051,Fe/O=0.055.a These enhanced abun-dances may be an arte-fact of SWCX contami-nation;see§7.2.shown in thoseﬁgures underpredict the neon and mag-nesium emission.We discuss these results in§§7.2and§7.5.In§7.2we discuss the possibility that the enhanced neon and mag-nesium emission is in fact due to SWCX contamination of these lines,rather than being due to these elements being enhanced in the halo.On the other hand,in§7.5 we discuss the implications of neon really being enhanced in the halo with respect to oxygen and iron.PARING THE SUZAKU AND XMM-NEWTONSPECTRAFor comparison,Table7also contains the results of the analysis of the XMM-Newton spectra from the same ob-servation directions by Henley et al.(2007).The Model9 results are taken directly from their“standard”model. However,it should be noted that Henley et al.(2007) used APEC to model all of their data,whereas in the analysis described above we used the Raymond-Smith code to model the ROSAT R1–3bands.We have therefore reanalyzed the XMM-Newton+ROSAT spec-tra,this time using the Raymond-Smith code to model the ROSAT R1–3bands,and using APEC to model the ROSAT R4–7bands and the XMM-Newton spec-tra.This new analysis allows a fairer comparison of our XMM-Newton results with our Suzaku results.The XMM-Newton spectra we analyzed are identical to those used by Henley et al.(2007)–see that paper for details of the data reduction.We added a broken power-law to the model to take into account soft-proton contamination in the XMM-Newton spectra.This broken power-law was not folded through XMM-Newton’s eﬀective area, and was allowed to diﬀer for the on-and oﬀ-ﬁlament datasets(see Henley et al.2007).The presence of this contamination means we cannot independently constrain the normalization of the extragalactic background.We therefore freeze the on-and oﬀ-ﬁlament normalizations at the Suzaku-determined values.The results of this new analysis are presented as Model10in Table7.As can be seen,there is poor agreement between the best-ﬁt parameters of the Suzaku +ROSAT model(Model2)and the XMM-Newton+ ROSAT model(Model10).We believe this discrepancy is due to an extra emission component in the XMM-Newton spectra.In Figure7we plot the diﬀerences between the XMM-Newton spectra and our best-ﬁtting0.10.20.30.40.5data-model(ctss-1keV-1)On filament MOS 1MOS 2 00.10.20.30.40.50.5125channel energy (keV)Off filament MOS 1MOS 2Fig.7.—The excesses in our on-ﬁlament(top)and oﬀ-ﬁlament (bottom)XMM-Newton spectra over our best-ﬁtting model to the Suzaku+ROSAT data(Table7,Model2).The gap in the data between1.4and1.9keV is where two bright instrumental lines have been removed.Suzaku+ROSAT model(Model2from Table7).To our best-ﬁtting Suzaku+ROSAT model we have added a broken power-law to model the soft-proton contami-nation in the XMM-Newton spectra.The parameters of this broken power-law are frozen at the values de-termined from theﬁtting to the XMM-Newton spectra described in the previous paragraph.The on-ﬁlament XMM-Newton spectra show excess line emission at∼0.57 and∼0.65keV,most likely due to O vii and O viii,re-spectively.The features in the oﬀ-ﬁlament spectra are not as clear.However,there appears to be excess O vii emission in the MOS1spectrum,and excess emission at ∼0.7keV(of uncertain origin)and∼0.9keV(proba-bly Ne ix)in the MOS2spectrum.We can estimate the signiﬁcance of the excess emission by calculatingχ2 for the XMM-Newton data compared with the Suzaku +ROSAT model.We concentrate on the excess oxy-gen emission and calculateχ2for the0.5–0.7keV energy range.Weﬁndχ2=106.28for24degrees of freedom for the on-ﬁlament spectra,andχ2=43.03for22degrees of freedom for the oﬀ-ﬁlament spectra.These correspond toχ2probabilities of2.5×10−12and0.0047,respectively, implying that the excesses are signiﬁcant in both sets of spectra at the1%level.We measure the intensities of the extra oxygen emis-sion byﬁttingδ-functions at E=0.570keV and 0.654keV to the excess spectra in Figure7.Weﬁt theseδ-functions simultaneously to the on-and oﬀ-ﬁlament XMM-Newton excess spectra.The intensities of the excess oxygen emission in the XMM-Newton spec-tra over the best-ﬁtting Suzaku+ROSAT model are 3.8±0.5L.U.(O vii)and1.4±0.3L.U.(O viii).We believe that this excess oxygen emission is due to SWCX contamination in our XMM-Newton spectra.As noted in the Introduction,this was not taken into ac-count in the original analysis of the XMM-Newton spec-tra.This is because the solar windﬂux was steady and slightly below average during the XMM-Newton obser-vations,leading Henley et al.(2007)to conclude that SWCX contamination was unlikely to be signiﬁcant.We discuss the SWCX contamination in our spectra further in§6.5.MEASURING THE OXYGEN LINES8HENLEY AND SHELTONAs well as using the above-described method to sep-arate the LB emission from the halo emission,we mea-sured the total intensities of the O vii complex and O viii line at∼0.57and∼0.65keV in each spectrum.These lines are a major component of the SXRB,accounting for the majority of the observed ROSAT R4diﬀuse back-ground that is not due to resolved extragalactic discrete sources(McCammon et al.2002),and are easily the most prominent lines in our Suzaku spectra.To measure the oxygen line intensities,we used a model consisting of an absorbed power-law,an absorbed APEC model whose oxygen abundance is set to zero,and two δ-functions to model the oxygen lines.As in the pre-vious section,the power-law models the extragalactic background,and its photon index was frozen at1.46 (Chen et al.1997).The APEC model,meanwhile,mod-els the line emission from elements other than oxygen, and the thermal continuum emission.The absorbing columns used were the same as those used in the ear-lier Suzaku analysis.As with our earlier analysis,we multiplied the whole model by a vphabs component to model the contamination on the optical blockingﬁlter which is in addition to that included in the CALDB con-tamination model(see§3.2).Weﬁx the carbon column density of this component at0.28×1018cm−2(Table7, Model2).Weﬁt this model simultaneously to our on-and oﬀ-ﬁlament spectra.However,all the model param-eters were independent for the two directions,except for the oxygen line energies–these were free parameters in theﬁt,but were constrained to be the same in the on-and oﬀ-ﬁlament spectra.We used essentially the same method to measure the oxygen line intensities in our XMM-Newton spectra.However,we did not use a vphabs contamination model,and,as before,we added a broken power-law to model the soft-proton contami-nation.Table3gives the energies and total observed intensities of the O vii and O viii emission measured from our Suzaku and XMM-Newton spectra.We can use the diﬀerence in the absorbing column for the on-and oﬀ-ﬁlament directions to decompose the ob-served line intensities into foreground(LB+SWCX)and background(halo)intensities.If I fg and I halo are the in-trinsic foreground and halo line intensities,respectively, then the observed on-ﬁlament intensity I on is given byI on=I fg+e−τon I halo,(1) whereτon is the on-ﬁlament optical depth at the energy of the line.There is a similar expression for the observed oﬀ-ﬁlament intensity I oﬀ,involving the oﬀ-ﬁlament op-tical depthτoﬀ.These expressions can be rearranged to giveI fg=eτon I on−eτoff I oﬀe−τon−e−τoff.(3) For the purposes of this decomposition,we use the Ba l uci´n ska-Church&McCammon(1992)cross-sections (with an updated He cross-section;Yan et al.1998)with the Wilms et al.(2000)interstellar abundances.We use the cross-sections at the measured energies of the lines. For the Suzaku measurements,the cross-sections we use are7.17×10−22cm2for O vii(E=0.564keV)and 4.66×10−22cm2for O viii(E=0.658keV).For the XMM-Newton O vii emission we use a cross-section of 7.03×10−22cm2(E=0.568keV).We cannot decom-pose the XMM-Newton O viii emission because the on-ﬁlament O viii line is brighter than the oﬀ-ﬁlament line. This gives rise to a negative halo intensity,which is un-physical.The results of this decomposition are presented in Ta-ble4.Note that the foreground oxygen intensities mea-sured from the Suzaku spectra are consistent with zero. This is consistent with our earlierﬁnding that the Suzaku spectra are consistent with there being no local emis-sion in the Suzaku band(see§3.2).The diﬀerence be-tween the foreground O vii intensities measured from our XMM-Newton and Suzaku spectra is5.1±3.1L.U.. This is consistent with the O vii intensity measured from the excess XMM-Newton emission over the best-ﬁtting Suzaku+ROSAT model(3.8±0.5L.U.;see§4).The halo O vii intensities measured from our XMM-Newton and Suzaku spectra are consistent with each other.This is as expected,as we would not expect the halo intensity to signiﬁcantly change in∼4years.6.MODELING THE SOLAR WIND CHARGE EXCHANGEEMISSIONIn§§4and5we presented evidence that our XMM-Newton spectra contain an extra emission component,in addition to the components needed to explain the Suzaku spectra.In particular,the O vii and O viii emission are enhanced in the XMM-Newton spectra.We attribute this extra component to SWCX emission,as it seems unlikely to be due to a change in the Local Bubble or halo emission.This extra component helps explain why our XMM-Newton and Suzaku analyses give such diﬀerent results in Table7.Previous observations of SWCX have found that in-creases in the SWCX emission are associated with en-hancements in the solar wind,as measured by ACE. These enhancements consist of an increase in the pro-tonﬂux,and may also include a shift in the ionization balance to higher ionization stages(Snowden et al.2004; Fujimoto et al.2007).In§6.1,we present a simple model for heliospheric and geocoronal SWCX emission,and use contemporaneous solar wind data from the ACE and WIND satellites to determine whether or not the ob-served enhancement of the oxygen lines in the XMM-Newton spectra is due to diﬀerences in the solar wind between our two sets of observations.In addition to the variability associated with solar wind enhancements,the heliospheric SWCX intensity is also expected to vary during the solar cycle,due to the dif-ferent states of the solar wind at solar maximum and solar minimum(Koutroumpa et al.2006).As our two sets of observations were taken∼4years apart,at diﬀer-ent points in the solar cycle,in§6.2we examine whether the SWCX intensity variation during the solar cycle can explain our observations.6.1.A Simple Model for Heliospheric and GeocoronalSWCX Emission6.1.1.The BasicsA SWCX line from a X+n ion of element X results from a charge exchange interaction between a X+(n+1) ion in the solar wind and a neutral atom.The intensity of that line therefore depends on the density of X+(n+1)。

First Light Measurements with the XMM-Newton Reflection Grating Spectrometers Evidence for

a r X i v :a s t r o -p h /0011018v 1 1 N o v 2000Astronomy &Astrophysics manuscript no.(will be inserted by hand later)First Light Measurements with the XMM-Newton Reﬂection Grating Spectrometers:Evidence for an Inverse First Ionisation Potential Effect and Anomalous Ne Abundance inthe Coronae of HR 1099A.C.Brinkman 1,E.Behar 2,M.G¨u del 3,M.Audard 3,A.J.F.den Boggende 1,G.Branduardi-Raymont 4,J.Cottam 2,C.Erd 5,J.W.den Herder 1,F.Jansen 5,J.S.Kaastra 1,S.M.Kahn 2,R.Mewe 1,F.B.S.Paerels 2,J.R.Peterson 2,A.P.Rasmussen 2,I.Sakelliou 4,and C.de Vries 11Space Research Organization of the Netherlands,Sorbonnelaan 2,3584CA Utrecht,The Netherlands2Columbia Astrophysics Laboratory,Columbia University,550West 120th Street,New York,NY 10027,USA 3Paul Scherrer Institute,W¨u renlingen and Villigen,CH-5235Switzerland4Mullard Space Science Laboratory,University College London,Dorking RH56NS,United Kingdom5Astrophysics Division,Space Science Department of ESA,ESTEC,2200AG Noordwijk,The NetherlandsReceived –;accepted –Abstract.The RS CVn binary system HR 1099was extensively observed by the XMM-Newton observatory in February 2000as its ﬁrst-light target.A total of 570ks of exposure time was accumulated with the Reﬂection Grating Spectrometers(RGS).The integrated X-ray spectrum between 5–35˚Ais of unprecedented quality and shows numerous features attributed to transitions of the elements C,N,O,Ne,Mg,Si,S,Fe,and Ni.We perform an in-depth study of the elemental composition of the average corona of this system,and ﬁnd that the elemental abundances strongly depend on the ﬁrst ionisation potential (FIP)of the elements.But different from the solar coronal case,we ﬁnd an inverse FIP effect ,i.e.,the abundances (relative to oxygen)increase with increasing FIP .Possible scenarios,e.g.,selective enrichment due to Ne-rich ﬂare-like events,are discussed.Key words.atomic processes —line:formation —nuclear reactions,nucleosynthesis,abundances —stars:coronae —stars:individual (HR 1099)—X-rays:stars1.IntroductionThe outer atmospheres of stars are of fundamental importance to the chemical enrichment of interstellar space through stel-lar winds and the ejection of energetic particles (Meyer 1985).Interestingly,the elemental composition of the hot magnetic so-lar corona and the solar wind are markedly different from the Sun’s surface composition,which suggests that some fraction-ation process in the denser layers of the solar atmosphere must selectively enrich the corona with particular elements.Various observations of the solar corona,the solar wind,solar energetic particles,and cosmic rays have established that elements with a ﬁrst ionisation potential (FIP)below 10eV are overabundant by factors of up to 4with respect to the solar photosphere,unlike elements with a high FIP that maintain their photospheric abun-dance (Feldman 1992,see also Meyer 1985).This is usually re-ferred to as the “FIP effect”.However,the physics of this frac-tionation mechanism is poorly parative abun-dance studies of stars other than the Sun may shed light on the underlying physics.2 A.C.Brinkman et al.:XMM-Newton RGS observations of HR1099Fig.1.First order spectrum of HR1099as measured by the RGS instruments onboard XMM-Newton.The most promi-nent spectral features are indicated justabove the spectrum by the correspond-ing emitting ion.present observation is HR1099,one of the X-ray brightest and most magnetically active stellar systems,which was observed early in the programme.HR1099has been the target of previ-ous EUVE and low-resolution X-ray spectroscopic investiga-tions(e.g.Pasquini et al.1989,Grifﬁth&Jordan1998).Ayres et al.(2001)qualitatively compare the Chandra HETG spec-tra of HR1099and Capella,the latter being a somewhat cooler and less active coronal source.A companion paper(Audard et al.2001)to the present work discusses spectral variability in HR1099,based on selected XMM-Newton RGS data sets. 2.Observations and AnalysisThe XMM-Newton X-ray observatory incorporates a pay-load with two identical high-resolution Reﬂection Grating Spectrometers(RGS;den Herder et al.2001)with a spectral resolution of about0.06˚A FWHM.The spectral band cov-ered by the RGS ranges from5to35˚A.This range contains strong lines of C,N,O,Ne,Mg,Si,S,Ca,Fe,and Ni.The clearly detected eight consecutive charge states of Fe are of crucial importance for the present analysis.HR1099was ob-served in January and February2000for a total of570ksec yielding approximately1.3million source counts.The spec-trum was extracted by counting events inside a narrow spatial window along the dispersed CCD image,and then by using the CCD energy resolution to separate the spectral orders.The ex-tracted totalﬁrst-order spectrum between7and35˚A is shown in Fig.1.3.ResultsIn optically thin coronal plasmas,the power of a particular spectral line depends on the collisional and radiative transition rates.In this work,these quantities are calculated by means of the relativistic Hebrew University Lawrence Livermore Atomic Code(Bar-Shalom et al.1998).Owing to the high spectral resolution,we can directly scale the contribution of each ion species to match the observation.The atomic-level models are each calculated for a single electron temperature,T max, at which the relevant ion has its maximum abundance.Fig.2 shows the theoretical spectrum from7–18˚A(red curve),which also includes a phenomenological continuum component(blue) corresponding to a bremsstrahlung spectrum(which dominates over any other continuum components at the temperatures in question),compared with the HR1099data points(black).As can be seen,the agreement is very good.The majority of lines in this region are due to highly ionised Fe as demonstrated by the individual ion curves at the bottom of theﬁgure.The con-tribution of each ion species yields a direct measurement of the amount of coronal gas in the conditions appropriate for that ion.With the distance d(=29pc)to HR1099,the calculated line power and the measured lineﬂux can be used to ob-tain the emission measure EM(in cm−3),which is a product of the electron density and the H-ion density,integrated over the emitting volume.We assume that each ion emits mostly around the temperature T max.This leads to aﬁrst approxi-mation of the emission measure distribution.This distribution is sufﬁciently accurate for our purposes if no steep gradients are present.Finally,in order to plot the EM,we use the ionic abundances from ionisation balance calculations(Arnaud& Raymond1992,Mazotta1998)and assume,as aﬁrst step,ele-mental abundances of the solar photosphere.The resulting EM as a function of the electron temperature T e is shown in Fig.3 for ions of C,N,O,Ne,Mg,S,Fe,and Ni.These ions sample the HR1099hot coronae in the T e range from kT e=100eV to1500eV and illustrate the monotonic increase of the EM in this range.The errors on the temperature for each ion represent the range of T e for which the ion fraction exceeds50%of its maximum value,except for the He-like ions O6+,Ne8+,and Mg10+,for which line ratio methods(Gabriel&Jordan1969) are used to determine T e more precisely and the bars reﬂect a 25%uncertainty.HR10993Fig.2.Observed HR1099ﬂux spec-trum(black dots)compared with thetotal calculated spectrum(red curve).The inferred thermal bremsstrahlungcontinuum is represented by the bluecurve.The eight separate plots at thebottom show the relative individual-ioncontributions of Fe16+to Fe23+(not toscale).The non-Fe lines(O,Ne,Mg,and Ni)are indicated explicitly.Table1.Abundance estimates a of the HR1099coronae,Solar photosphere b and Solar corona c relative to ONe21.560.570.150.15 1.0 3.8N14.530.170.120.13 1.08 1.4O13.61 1.0 1.0 1.0 1.0 1.0C11.260.280.470.50 1.060.59S10.360.00850.0190.024 1.260.45Fe7.870.00940.0380.16 4.210.25Mg7.640.0430.0450.18 4.00.95Ni7.630.000700.00210.0089 4.240.334 A.C.Brinkman et al.:XMM-Newton RGS observations of HR 10993.The empirical emission measure O 6+,O 7+,Ne 8+,Ne 9+,Mg 10+,11+,Fe 16+through Fe 23+,Ni 18+,5+,N 6+and S 13+(in order from left right for each element).The straight is the best-ﬁt power-law for the Fealone.Fig.4.Enrichment of elemental abun-dances in the HR 1099coronae rela-tive to (solar)photospheric values,plot-ted as a function of the ﬁrst ionisation potential (FIP).All values are given error bars of 20%.Note the increase with FIP,exactly opposite the trend ob-served in the average solar corona.the respective ratios of the solar corona,the contrast is even sharper.In fact,HR 1099reveals systematically a reversed FIP effect,which is illustrated in Fig.4;the coronal enrichment in-creases monotonically with the FIP.The only exception may be Mg (low FIP).Ne is strongly enhanced,by a factor of 3.8compared to the solar photosphere.This enhancement is rem-iniscent of a similar anomaly observed in a subset of solar ﬂares (Schmelz 1993,Murphy et al.1991).Interestingly,ﬂares are thought to be of great importance to the heating of ap-parently quiet coronae of active stars (Doyle &Butler 1985,G¨u del 1994,Audard et al.2000)and have been suggested to be pivotal for the elemental fractionation process as well (Wang1996,G¨u del et al.1999,Mewe et al.1997).Continuous mix-ing of photospheric and coronal material in low-lying loops through frequent ﬂares could suppress the solar-like FIP effect (Feldman &Widing 1990)in these stars altogether.It is generally believed that the fractionation is the result of diffusive processes across magnetic ﬁeld lines somewhere in the chromosphere,where the ionisation fraction differs sub-stantially among elements of high or low FIP (von Steiger &Geiss 1989).Low FIP elements can then be accelerated into the quiescent corona,leaving an excess of high FIP elements in the chromosphere (Schmelz 1993).If this chromospheric layer can be accessed by ﬂares (Schmelz 1993),then a high-FIP-A.C.Brinkman et al.:XMM-Newton RGS observations of HR10995rich plasma may be heated and lifted into dense coronal X-ray loops giving rise to the reversed coronal FIP effect observed in HR1099.In any case,it is clear that the simple mechanism in-volving steady-state diffusion,which has been proposed for the quiet solar corona,does not operate in more active stars such as HR1099.Since cosmic rays do show a low FIP enhancement, it is unlikely that such active stars can be major contributors to the cosmic rayﬂux.The wide survey of stellar coronae planned with the RGS on XMM-Newton will provide more general an-swers to the questions raised by these observations. Acknowledgements.We thank the many teams in industry and at the scientiﬁc institutions and ESA,for building such an excellent space-craft and payload-and operating it.We are particularly grateful to Mr.Robert Lain´e and his project team at ESA for leading the project.We thank the XMM-Newton SOC staff for making the com-missioning phase data available to us.SRON isﬁnancially supported by the Netherlands Organization for Scientiﬁc Research(NWO). The Columbia University team acknowledges generous support from the National Aeronautics and Space Administration.The PSI group is supported by the Swiss Academy of Natural Sciences and the Swiss National Science Foundation(grants2100-049343and2000-058827).MSSL acknowledges support from the Particle Physics and Astronomy Research Council.This work is based on observations ob-tained with XMM-Newton,an ESA science mission with instruments and contributions directly funded by ESA member states and the USA (NASA).ReferencesAntunes,A.,Nagase,F.,White,N.E.1994,ApJ,436,L83 Arnaud,M.,Raymond,J.1992,ApJ,398,394Audard,M.,G¨u del,M.,Drake,J.J.,Kashyap,V.2000,ApJ,in press Audard,M.,G¨u del,M.,Mewe,R.2001,A&A,365(this issue) Ayres,T.R.,Brown,A.,Osten,R.A.,Huenemoerder,D.P.,Drake,J.J., Brickhouse,N.S.,Linsky,J.L.2001,ApJ,in pressBar-Shalom,A.,Klapisch,M.,Goldstein,W.H.,Oreg,J.1998,The HULLAC Code for Atomic Physics,unpublishedden Herder,J.W.,Brinkman,A.C.,Kahn,S.M.,et al.2001,A&A,365 (this issue)Doyle,J.G.,Butler,J.1985,Nature,318,378Drake,J.J.,Laming,J.M.,Widing,K.G.1995,ApJ,443,393 Drake,J.J.,Laming,J.M.,Widing,K.G.1997,ApJ,478,403 Drake,S.A.1996,in:X-ray Measurements of Coronal Abundances, 6th Ann.Astrophysics Conference,ASP Conf.Proc.,S.S.Holt,G.Sonneborn(eds.),San Francisco:ASP,p.215Feldman,U.1992,Phys.Scripta46,202Feldman,U.,Mandelbaum,P.,Seely,J.F.,Doschek,G.A.,Gursky,H.1992,ApJS,81,387Feldman,U.,Widing,K.G.1990,ApJ,363,292Gabriel,A.H.,Jordan,C.1969,MNRAS,145,241G¨u del,M.1994,ApJS,90,743G¨u del,M.,Linsky,J.L.,Brown,A.,Nagase,F.1999,ApJ,511,405 H´e noux J.-C.1995,Adv.Space Res.,15,23Mazzotta,P.,Mazzitelli,G.,Colafrancesco,S.,Vittorio,N.1998, A&ASS,133,403Meyer,J.P.1985,ApJ,57,173Mewe,R.,Kaastra,J.S.,van den Oord,G.H.J.,Vink,J.,Tawara,Y.1997,A&A,320,147Murphy,R.J.,Ramaty,R.,Kozlovsky,B.,Reames,D.V.1991,ApJ, 371,793Pasquini,L.,Schmitt,J.H.M.M.,Pallavicini,R.1989,A&A,226,225Randich,S.,Giampapa,M.S.,Pallavicini,R.1994,A&A,283,893 Schmelz,J.T.1993,ApJ,408,373von Steiger,R.,Geiss,J.1989,A&A,225,222Wang,Y.-M.1996,ApJ,464,L91White,N.E.1996,in:Cool stars,Stellar Systems and the Sun,ASP Conf.Series,109,p.193。

脉冲星测量技术用于深空探测器自主导航的原理及方法

obs4376延迟量的确定由上述式并结合延迟量所受到的时间和空间效应有观测方程其中是飞行器的探测器在初始时刻观测到的第一个脉冲星信号与当前脉冲到达时间之差是太阳系质心相对于太阳质心的位置矢量urur4476延迟量的确定上式中的第一项称为一阶多普勒延迟即飞行器的位置矢量在脉冲星方向上的投影第二项表示周年视差影响第一项和第二项被合称为roemer迟该延迟是构成脉冲星信号观测量的主要因素
脉冲星测量技术用于深空探测器自主导航的原理及方法
1/76
脉冲星的发现
2/76
脉冲星机制
当一颗恒星变成超新星时，经过激烈变化后，留下满天膨胀的气体和微小物质，余下的核心直径只有几十到十几公里。超新星的内爆非常强烈，恒星原子里的质子和电子被紧紧地压缩在一起，抵消了它们的电荷，形成中子。这种中子星可以达到水密度的1014倍，有着极强的磁场，可以非常快速地旋转。因为磁轴不与旋转轴重合，二者一般具有一定的夹角，当脉冲星高速旋转时，辐射束将沿着磁场两极方向被抛出，随着脉冲星的自转，该辐射束周期性扫过探测器的视界，形成脉冲。
11/76
脉冲星导航技术研究历程
12/76
研究历程
a) 脉冲星导航思想的萌芽阶段 1. 脉冲星导航思想最早于 20 世纪 70 年代提出。 2. 1971 年，Reichley，Downs 和Morris 首次描述了把射电脉冲星作为时钟的思想。 3. 1974 年，Downs 在文献《Interplanetary Navigation Using Pulsation Radio Source》中提出一种基于射电脉冲星信号进行行星际导航的思想，标志着脉冲星导航思想的初步形成。但由于脉冲星的射电信号强度较弱，宇宙中的射电信号噪声强度大，导航中需要至少 25m 口径的天线接收信号，因此该方法很难在工程中实现。 4. 20 世纪 70 年代后期，天文观测在 X 射线波段能量范围 1~20keV、频率范围 2.5×1017 ~ 4.8×1017Hz的进展，促进了对 X 射线脉冲星特性的研究。 5. 1980 年 Downs 和 Reichley 提出测量脉冲星脉冲到达时间的技术。 1981 年 Chester 和Butman 在国际上第一次正式提出利用 X 射线脉冲星进行航天器导航的思想。

MCG-6-30-15的铁Kα线的研究

1.3 能谱拟合
我们用XSPEC软件拟合观测能谱。为了应用Ⅹ2检验，我们首先使能谱的每一个能量 bin里至少包含有20个光子。以下的拟合包含了银河系的吸收（各个源在视线方向上的中性氢柱密度可由 Heasoft软件包中全天中性氢分布得到［8］），还加上了源本征的吸收。
射电宁静AGNs的X射线能谱通常包含一个或多个连续谱成分以及一些吸收和发射的特征。MCG－6－30－15也同样如此。由于软 X射线（0．3－2ke V）波段的能谱比较复杂，我们首先拟合2－10 Ke V波段的能谱，通常一个简单的幂律谱能较好地拟合此能段的数据，如图1所示。从图1可看出，在6—7 Ke V能段有明显的的“线状”的剩余，即铁Kα线。拟合得到的参数分别为：幂律谱指数 T＝1．84±0．04，Reduced chi－squared（Ⅹ2）＝1．72（大于1．2，拟合不能接受）。当把幂律模型外推到0．3－2．0Ke V能段时，我们发现在软X射线能段有很强的超出，就是我们通常说的“软X射线超”或“软X射线过剩（soft X－ray excess）”，如图2所示。我们用一个黑体成分来代表“软X射线超”。拟合结果如图3所示，从图上可以看出，在软X射线波段，有明显的吸收。拟合得到黑体温度k T＝0．09±0．02 Ke V，幂律谱指数T＝1．69±0．02，Reduced chi－squared（Ⅹ2）＝1．30。接下来的拟合中，我们加了几条吸收边以及一个高斯成分来代表铁Kα线。拟合结果如图4所示。从图4可看出，此拟合效果较好。拟合得到幂律谱指数T＝1．85±0．04，黑体温度k T＝0．13±0．005Ke V，铁线的线心能量（LineE）为6．37±0．02Ke V，线宽（Sig ma）为 0．18±0．02 Ke V。Reduced chi－squared（Ⅹ2）＝1．15。

The XMM-Newton view of radio-quiet and X-ray dim isolated neutron stars

a r X i v :a s t r o -p h /0401075v 1 7 J a n 2004Mem.S.A.It.Vol.,1cSAIt 2002MemoriedellaF.HaberlMax-Planck-Institut f¨u r extraterrestrische Physik,Giessenbachstraße,85748Garching,GermanyAbstract.As part of the guaranteed time,guest observer and calibration pro-grams,XMM-Newton extensively observed a group of six thermally emitting,isolated neutron stars which are neither connected with a supernova remnant nor show pulsed radio emission.The high statistical quality of the EPIC data allows a detailed and homogeneous analysis of their temporal and spectral properties.Four of the six sources are now well established as X-ray pulsars:The 11.37s period discovered in EPIC-pn data of RX J0806.4−4123was conﬁrmed in a second XMM-Newton observation.In the case of the X-ray faintest of the six stars,RX J0420.0−5022the period marginally indicated in ROSAT data was not seen in the EPIC data,instead a 3.45s pulse period was clearly detected.Spectral variations with pulse phase were discovered for the known 8.39s pul-sar RX J0720.4−3125and also RBS1223.For the latter EPIC data revealed a double-peaked pulse proﬁle for a neutron star spin period of 10.31s.The X-ray continuum spectra of all six objects are consistent with a Planckian energy distribution with black-body temperatures kT between 40eV and 100eV.EPIC data of the pulsars RBS1223and RX J0720.4−3125revealed a broad absorption feature in their spectra at energies of 100-300eV and ∼260eV,re-spectively.The depth of this feature varies with pulse phase and may be caused by cyclotron resonance scattering of protons or heavy ions in a strong magnetic ﬁeld.In such a picture the inferred ﬁeld strength exceeds 1013Gauss,in the case of RX J0720.4−3125consistent with the value estimated from its pulse pe-riod derivative.A similar absorption feature in the RGS and EPIC spectra of RX J1605.3+3249was reported recently.Key words.X-rays:stars –stars:neutron –stars:magnetic ﬁelds1.IntroductionPresently more than half a dozen ROSAT-discovered X-ray dim isolated neutron stars (XDINs,for recent reviews see Treves et al.2 F.Haberl:Radio-quiet and X-ray dim isolated neutronstarsFig.1.EPIC-pn light curves folded on the pulse period for the four pulsars among the XDINs.remnants.Some of them exhibit pulsationsin their X-rayﬂux indicating the neutronstar rotation period.Although there is littledoubt that the soft X-rays are of thermalorigin from the surface of an isolated neu-tron star,the details for the formation ofthe spectrum are not clear.XMM-Newtonobserved six XDINs as part of the guaran-teed time,guest observer and calibrationprograms.Here I summarize some of theﬁrst results of a homogeneous analysis oftheir temporal and spectral properties.2.X-ray pulsationsThe second brightest of the known XDINsis RX J0720.4−3125and was theﬁrstdiscovered as X-ray pulsar(Haberl et al.1997)with a period of8.39s.XMM-Newton observed RX J0720.4−3125sixtimes.As example the folded EPIC-pnlight curve from satellite revolution534is drawn in Fig.1which shows a pulsedfraction of about11%.In Chandra dataof RBS1223=1RXS J130848.6+212708Hambaryan et al.(2002)discovered pulsa-tions with a period of5.16s and XMM-Newton observations revealed that the trueneutron star spin period is more likely10.31s.This is supported by pulse phase de-pendent hardness ratio variations whichare diﬀerent for the two intensity max-ima(Haberl et al.2003).RBS1223ex-hibits the deepest modulation of the knownXDINs in its double peaked pulse proﬁleof18%(Fig.1).Theﬁrst XMM-Newtonobservation of RX J0806.4−4123revealeda6%modulation(Fig.1)with a periodof11.37s(Haberl&Zavlin2002)whichwas conﬁrmed in a second observation(Haberl et al.2004a).Four XMM-Newtonobservations of RX J0420.0−5022did notconﬁrm the22.7s pulsations originally in-dicated in ROSAT data,but clearly re-F.Haberl:Radio-quiet and X-ray dim isolated neutron stars3veal a3.45s period.The pulsed fraction is about12%(Fig.1).In Table1the de-rived pulse periods for the four pulsars are summarized together with pulse period derivatives assuming linear period changes between multiple XMM-Newton observa-tions.In most cases the time base line is too short to determine precise˙P values, only for RX J0720.4−3125observations by diﬀerent satellites cover already more than 10years(Zane et al.2002;Kaplan et al. 2002).3.X-ray spectraThe X-ray spectra of XDINs obtained by the ROSAT PSPC were all consistent with black-body emission little attenuated by interstellar absorption suggesting that the objects are close-by.To look for ab-sorption features which may be created by heavy chemical elements in the stel-lar atmosphere high resolution spectra of RX J1856.4−3754were obtained by the LETGS on Chandra(Burwitz et al.2001, 2003).Surprisingly,no signiﬁcant narrow features were detected.A spectral analysis(in a homogenous way using the latest calibration data)of the EPIC-pn spectra,which are of unprece-dented statistical quality,shows that in sev-eral cases a black-body model yields un-satisfactoryﬁts.The strongest deviations are seen from RBS1223and Haberl et al. (2003)demonstrate that a non-magnetic atmosphere models can neither explain the observed spectrum.However,it was found that adding an absorption feature mod-eled by a broad(100eV)Gaussian line at an energy between100and300eV to the Planckian continuum yields accept-ableﬁts.Similarly,but at a higher en-ergy of450eV and therefore inside the sensitive band of the RGS instruments van Kerkwijk(2004)presented the detec-tion of a broad absorption feature in the spectra of RX J1605.3+3249.Haberl et al. (2004b)report a broad absorption fea-ture in the EPIC-pn spectra of the pul-sar RX J0720.4−3125at an energy of270eV andﬁnd that the depth of the feature varies with pulse phase by a factor of∼2. Finally,also spectra of RX J0420.0−5022 indicate a possible absorption line at330 eV(Haberl et al.2004a).In Table1the spectral parameters in-ferred from theﬁts using the two mod-els(absorbed black-body and absorbed black-body with Gaussian-shaped absorp-tion line)are listed.Theﬁrst value given for column density N H and black-body tem-perature kT refers to the model without ab-sorption line and the second to the model with line added.4.DiscussionFour of the six X-ray-dim isolated neu-tron stars are now known as X-ray pul-sars with spin periods between3.45s and 11.37s.The fraction of pulsedﬂux in their folded X-ray light curves ranges between 6%and18%.RBS1223,RX J0720.4−3125 and RX J0420.0−5022show hardness ratio variations with pulse phase(Haberl et al. 2003,2004b,a)while for RX J0806.4−4123 the shallow modulation makes the signiﬁ-cant detection of such an eﬀect more diﬃ-cult(Haberl et al.2004a).Pulse phase re-solved spectra for theﬁrst three pulsars show that the temperature changes only marginally with pulse phase.Any model for the pulsed emission from this group of iso-lated neutron stars must be able to explain this behaviour.An important piece of information comes from the detection of broad ab-sorption features in the X-ray spectra of several XDINs.At least in the case of RX J0720.4−3125the depth of the feature varies strongly with pulse phase.A likely interpretation of these features is cyclotron resonance absorption which can be ex-pected in spectra from magnetized neutron stars withﬁeld strength B in the range of 1010–1011G or2×1013–2×1014G if caused by electrons or protons,respectively.In the case of RX J0720.4−3125the measured˙P, if interpreted as magnetic dipole braking, rules out the lower range for B,leaving pro-4 F.Haberl:Radio-quiet and X-ray dim isolated neutron starsTable1.X-ray-dim isolated neutron stars observed by XMM-NewtonRX J0420.0−5022 3.453<9×10−12 1.0/2.045/45330?6? RX J0720.4−31258.391(3−6)×10−14 1.4/0.885/842705 RX J0806.4−412311.371<2×10−120.496−RBS122310.313<6×10−127.1/4.195/86100-3002-6 RX J1605.3+3249−−0.3/0.996/934509 RX J1856.5−3754−−0.960−1?F.Haberl:Radio-quiet and X-ray dim isolated neutron stars5 2002,MNRAS,334,345。

What can XMM-Newton tell us about the spin periods of Intermediate Polars

a r Xiv:as tr o-ph/49638v223Nov24The astrophysics of cataclysmic variables and related objects ASP Conference Series,Vol.?,2004J.M.Hameury &sota What can XMM-Newton tell us about the spin periods of Intermediate Polars?P.A.Evans &Coel Hellier Astrophysics Group,School of Chemistry and Physics,Keele University,Keele,Staﬀordshire,ST55BG Abstract.XMM-Newton ’s unprecedented combination of spectral resolution and high throughput allows us to perform the best phase-resolved X-ray analysis of intermediate polars to date.The Optical Monitor gives optical/UV photom-etry simultaneously with the X-ray data.We present a comprehensive study of X-ray spin pulses in IPs,giving spin-pulses and hardness ratios for every IP looked at with XMM-Newton to date.1.Introduction In an intermediate polar (IP)the white dwarf has a magnetic ﬁeld strong enough to aﬀect the accretion ﬂow,but not strong enough to synchronise the white dwarf rotation with the binary period;the magnetic dipole is also inclined to the spin axis of the star.Generally accretion occurs from a magnetically truncated accretion disc;at some radius the disc material threads to the white dwarf’s ﬁeld lines and is channeled towards the gravitationally preferable pole in large accretion curtains.Stand-oﬀshocks form above each pole,resulting in hard X-ray emission.As the white dwarf rotates,our view of these regions changes and the position of the accretion curtains relative to our line of sight also varies,giving rise to spin-period modulations (e.g.Hellier,Cropper,&Mason 1991).To date eleven IPs have been observed with XMM-Newton .In this paper we present a compilation of X-ray spin-pulse proﬁles,hardness ratios,and UV data from XMM ’s Optical Monitor (OM)for these systems.2.AO Piscium and V1223SagittariiAO Psc and V1223Sgr show sinusoidal X-ray and UV modulations (Fig.1).The fact that their spectra get harder at pulse minimum suggests that the modulation is caused by absorption.These are thus good examples of the accretion curtain model for X-ray modulation in IPs (e.g.Hellier et al.1991).3.HT CamelopardarisHT Cam’s X-ray spin-pulse (Fig.2)appears sinusoidal with a ﬂattened maxi-mum,and,judged from the softness ratio,appears to be energy independent.No pulse is detected in the UV,although the B -band shows sinusoidal variation.For more information,see de Martino et al.(2004b).12Evans and HellierFigure1.The spin pulses of AO Psc(left)and V1223Sgr(right).We showthe0.2–12keV X-ray data,the(0.2–4)/(6–12)keV softness ratio and(whereavailable)the UV(2050–2450˚A)data.4.RX J1548.2−4528The recently discovered IP RX J1548.2−4528shows a sinusoidal X-ray spin-pulse(Fig.2).Haberl,Motch&Zickgraf(2002)analysed this observation,and suggested that while there may be a small change in absorption with spin phase, the predominant source of modulation is a change in the visibility of the emission.5.EX HydraeEX Hya’s softness ratio has the same shape and phasing as the lightcurve (Fig.3),making it tempting to interpret this modulation as changing absorption in the accretion curtains(indeed,it was for this system that accretion curtains wereﬁrst proposed as the cause of X-ray modulation;Rosen,Mason,&C´o rdova 1988).However,various X-ray studies(e.g.Allan,Hellier,&Beardmore1998) have suggested that the upper emitting pole is periodically occulted by the white dwarf,giving rise to the modulation.Since such occultation will aﬀect the lower, cooler parts of the accretion column more than the higher,hotter ones,we see a deeper modulation at lower energies.6.AE AquariiAE Aqr diﬀers from many IPs as it is a rapid rotator(P spin∼33s),and thus its magneticﬁeld might be expelling material from the system like a propeller(e.g., Eracleous&Horne1996).It is unclear whether the X-ray emission arises fromXMM-Newton IP spin pulses3Figure2.As for Fig.1,but for HT Cam(left)and RX J1548.2−4528(right). The UV band for HT Cam is1800–2250˚A,and the B band(3900–4900˚A)is also shown.Figure3.As Fig.1,but for EX Hya(left)and AE Aqr(right).4Evans and HellierFigure4.Left:As for Fig.1,but for GK Per.Right:Spin-pulse proﬁle ofV405Aur.The upper panel shows the0.2–0.7keV band and the lower panelshows the0.7–12keV band.accretion,as is usual in IPs,or further out in the magnetosphere,as suggested by Ikhsanov(2001).However,the XMM-Newton data shows that the pulse proﬁle is sinusoidal and largely independent of energy.7.GK Perseii(in outburst)Watson,King,&Osborne(1985)suggested that the X-ray spin-pulse of GK Per in outburst was caused by increased absorption at spin minimum.Ishida et al.(1992)blamed changing absorption for the proﬁle in quiescence as well,but suggested that the outburst proﬁle may show evidence for occultation of the upper pole.Fig.4shows that the outburst softness ratio follows the lightcurve, making it more likely that absorption changes are responsible for this pulse,as Watson et al.(1985)claimed.Hellier,Harmer,&Beardmore(2004)agree with this,in analysis of an RXTE observation.8.V405AurigaeV405Aur diﬀers from IPs such as AO Psc,ﬁrst because it shows a soft blackbody component to its X-ray emission(Haberl et al.1994),and second because the blackbody is double-peaked on the spin period,while the harder emission is single-peaked but sawtoothed(Fig.4).Furthermore,de Martino et al.(2004a) and Evans&Hellier(2004)have shown that the absorption does not change with spin phase.Evans&Hellier(2004)suggest that if the angle between the magnetic and spin axes in this system is very high,then the double-peaked proﬁle of the blackbody emission is explained.XMM-Newton IP spin pulses5Figure5.Left:as Fig.1,but for FO Aqr.Right:PQ Gem spin folds forthe blackbody(upper panel),hard X-ray(lower panel)and UV(2450–3200˚A)components.9.FO AquariiThe X-ray pulse proﬁle of FO Aqr(Fig.5)shows X-ray minimum to occur after UV minimum.Many authors(e.g.Hellier1993;Beardmore et al.1998) have identiﬁed the‘notch’at phase0.7as the result of occultation of the upper accretion column,and the dip around phase0.5as arising from the accretion curtains intercepting our line of sight.Evans et al.(2004)note that,if this is the case,the upper pole will be pointed towards the observer nearly a quarter of a cycle before the accretion curtain dip.They thus suggest that the accretion curtains are twisted,explaining the lightcurve and softness ratio.10.PQ GeminorumPQ Gem shows a soft blackbody component as well as the hard,optically thin emission characteristic of IPs(Mason et al.1992).The spin-pulse proﬁle(Fig.5) diﬀers greatly from that of AO Psc,with UV maximum coinciding with X-ray minimum.Potter et al.(1997)and Mason(1997)have suggested that the blackbody modulation is caused by changing views of the accretion region as the white dwarf rotates.Maximum(phase0)occurs when the upper pole is towards us.Since this occurs after the dip around phase0.8,interpreted as absorption by the accretion curtains,Potter et al.(1997)and Mason(1997)suggest that PQ Gem accretes preferentially alongﬁeld lines preceding the magnetic pole.6Evans and HellierFigure6.Power spectrum of V2400Oph.The spin(ω)and beat(ω−Ω)frequencies are marked.11.V2400OphiuchiV2400Oph is thought to be the only discless IP,since it shows a dominant X-ray modulation on the spin-orbit beat period(Buckley et al.1995;Hellier& Beardmore2002).We have recently discovered evidence for a spin pulse during an XMM-Newton observation(Fig.6),theﬁrst in the X-ray band.ReferencesAllan,A.,Hellier,C.,&Beardmore,A.P.1998,MNRAS,295,167Beardmore,A.P.,Mukai,K.,Norton,A.J.,Osborne,J.P.,Hellier,C.1998,MNRAS, 297,337Buckley, D.A.H.,Sekiguchi,K.,Motch, C.,O’Donoghue, D.,Chen, A.-L., Schwarzenberg-Czerny,A.,Pietsch,W.,Harrop-Allin,M.K.1995,MNRAS,275, 1028Eracleous,M.,&Horne,K.1996,ApJ,471,427Evans,P.A.,Hellier,C.,Ramsay,G.,Cropper,M.2004,MNRAS,349,715Evans,P.A.,&Hellier,C.2004,MNRAS,353,447Haberl,F.,Thorstensen,J.R.,Motch,C.,Schwarzenberg-Czerny,A.,Pakull,M.,Sham-brook,A.,Pietsch,W.1994,A&A,291,171Haberl,F.,Motch,C.,&Zickgraf,F.-J.2002,A&A,387,201Hellier,C.,Cropper,M.,&Mason,K.O.1991,MNRAS,248,233Hellier,C.,&Beardmore,A.P.,2002,MNRAS,331,407Hellier,C.,Harmer,S.,&Beardmore,A.P.2004,MNRAS,349,710Ikhsanov,N.R.2001,A&A,374,1030Ishida,M.,Sakao,T.,Makishima,K.,Ohashi,T.,Watson,M.G.,Norton,A.J.,Kawada, M.,Koyama,K.1992,MNRAS,254,647de Martino,D.,Matt,G.,Belloni,T.,Haberl,F.,Mukai,K.2004a,A&A,415,1009 de Martino,D.,Matt,G.,Mukai,K.,Bonnet-Bidaud,J.M.,G¨a nsicke,B.T.,Haberl,F.,Mouchet,M.,2004b,in Hameury,J.M.,&Lasota,J.P.,eds.,The Physics ofCataclysmic Variables and Related Objects,ASP Conf.Ser,in pressMason,K.O.,et al.1992,MNRAS,258,749Mason,K.O.1997,MNRAS,285,493Potter,S.B.,Cropper,M.,Mason,K.O.,Hough,J.H.,Bailey,J.A.1997,MNRAS,285, 82Rosen,S.R.,Mason,K.O.,&C´o rdova,F.A.1988,MNRAS,231,549Watson,M.G.,King,A.R.,&Osborne,J.P.,1985,MNRAS,212,917。

欧空局对地观测卫星介绍综述

6 比利时Redu地面站
Redu地面站拥有多个地面终端，位于比利时的Ardennes区，距Redu村庄约1公里，在比利时卢森堡省内。终端提供C波段、L波段、S波段、 Ku波段和Ka波段的跟踪数据，并提供通信卫星在轨检测（in-orbit tests，IOT）。目前S波段站用于Integral，该站位于东经5.146 231度、北纬 50.002 685度，参考点位于海拔378.3 m。地球同步可见仰角超过10度，在54.9度和65.1度之间。该站装备有自己的紧急发电机系统，并通过完全重复的光纤线路与比利时通讯网络相连接。该站全天24 小时运行。站点19公顷的面积上装备着一套自动安全系统，还有相当大的空间做进一步扩充。它的位置远离RF干涉，站点还有所有必需的频率清除用 2019/4/7 于运转S、C、Ku和Ka频率。
2019/4/7
15 m天线位于北纬5.25 143 694度、西经52.80 466 242度，根据WGS-84地球椭圆体参数天线位于14.561 m。 Kourou技术设备由S波段和X波段卫星上/下链设备组成，包括一套自动跟踪设备、一套测距系统、一套频率和定时系统、一套监控系统和一套通讯系统。配有不间断电源设备。 Kourou站对ESA的卫星活动起着至关重要的作用，特别是在LEOP阶段。ESA要求地面站寻找一种灵活、稳健和低本高效的方法进行设施的运行和维护（包括内部和外部的通讯连接）。库鲁地面站是位于全球8个不同地方的11个站点所组成的网络站点之一，其主要作用是为航天器的跟踪、遥测和指挥（Tracking, Telemetry and Command，TT&C）提供支持。
2019/4/7
9 ESTRACK的合作站点
除欧空局自己的地面站外，ESOC还与其他组织有着地面站服务的合作协议，必要时还包括ESA的设备。重要的合作有： Malindi（MAL），位于肯尼亚，与罗马大学（意大利）合作； Santiago（AGO），位于智利，与智利大学空间研究中心合作； Svalbard（SVA），位于

基于电子束离子阱的天体X射线和极紫外辐射研究进展

基于电子束离子阱的天体X射线和极紫外辐射研究进展彭吉敏;仲佳勇;梁贵云;刘畅【摘要】观测光谱是研究天体的基础资料,对这些光谱的理解主要依靠理论模型.由于各模型依赖于众多的基本原子参数,且不同的模型对同一特定对象有不同的理解,所以对这些模型进行实验校准十分重要.电子束离子阱(electron beam ion trap,简称EBIT)能够产生与天体环境类似的等离子体环境,因此基于电子束离子阱的实验室光谱测量极大地促进了对天体X射线和极紫外辐射的研究.主要对以下内容进行介绍:卫星观测光谱和理论数值计算的现状及所面临的问题;基于电子束离子阱的类氖铁(Fe XVII)3C/3D重要诊断谱线比的最新研究进展;近年来实验室电子束离子阱对X射线和极紫外辐射的测量进展.最后,对现有的实验室光谱研究工作予以简要的展望.%Observed spectram is the basic data for celestial research. In recent years, with the enhancement of space, equipment, imaging technologies and simulation techniques, spec-trometers with larger effective area have provided abundant X-ray and extreme ultra-violet spectra with high resolution. These spectral data which will post deep insights on different ob jects, such as plasma heating, accretion of compact ob jects, energy storage and release, emission structure, feedback of supernova, as well as kinetics and dynamics in X-ray/EUV emitters, are mainly dissected by theoretical calculations. However, the models depend on many elementary atomic parameters and there are different understanding for the same ob ject between different models, it is necessary to calibrate these models by laboratory measurement. The electron beam ion trap which is an inexpensive compact device is an ideal facility for spectroscopicbenchmark due to its overlapping electron density and temperature with as-trophysical plasma in stellar corona, supernova remnants and so on. Therefore, the spectral measurement based on the electron beam ion trap greatly promotes the study of X-ray and extreme ultraviolet radiation. This article will focus on the following: current situation of spectra of satellite observations and theoretical calculation and the problems; the latest re-search progress of the important diagnostic line ratio(Fe XVII's 3C/3D) based on EBIT;recent advances in laboratory measurement of X-rays and extreme ultraviolet radiation based on electron beam ion trap. Finally, a brief prospect for the laboratory spectroscopy research will be given.【期刊名称】《天文学进展》【年(卷),期】2017(035)004【总页数】12页(P417-428)【关键词】电子束离子阱;X射线;极紫外辐射;原子数据;高分辨率【作者】彭吉敏;仲佳勇;梁贵云;刘畅【作者单位】北京师范大学天文系,北京 100875;中国科学院光学天文重点实验室国家天文台,北京 100012;北京师范大学天文系,北京 100875;中国科学院光学天文重点实验室国家天文台,北京 100012;北京师范大学天文系,北京 100875【正文语种】中文【中图分类】P141了解天体辐射与光谱信息依赖于天体辐射成像和光谱观测，光谱获得主要基于高分辨率光谱仪观测。

Merging clusters of galaxies observed with XMM-Newton

a r X i v :a s t r o -p h /0501377v 1 18 J a n 2005Merging clusters of galaxies observed withXMM-NewtonE.Belsole a J-L.Sauvageot b G.W.Pratt c and H.Bourdin da H.H.Wills Physics Laboratory -University of Bristol,Tyndall Avenue,Bristol BS81TL,UKb Service d’Astrophysique,CE-Saclay,L’orme des Merisiers,91191Gif sur Yvette Cedex,Francec MPE Garching,Giessenbachstraße,85748Garching,Germanyd Dipartimentodi Fisica,Universit`a degli studi di Roma Tor Vergata,Via della Ricerca Scientiﬁca 1,00133Roma,Italy1IntroductionAs the largest assembled structures in the Universe,clusters of galaxies are commonly thought to form by gradual accretion of matter alongﬁlaments and by interaction and merging with previously formed structures.Hydrody-namical simulations(e.g.Rowley et al.2004and references therein)predict that merger events strongly aﬀect the physical characteristics of the intra-cluster medium(ICM).In particular the temperature structure is thought to be an excellent indicator of the cluster dynamical state and formation his-tory.Spectro-imaging observations with XMM-Newton and Chandra allow the building of precise temperature maps,enabling deeper investigation of the dynamical processes of cluster formation and the eﬀect of the mergers on the ICM(e.g.,Markevitch et al.,2000;Neumann et al.,2003;Henry et al.,2004; Krivonos et al.,2003;Vikhlinin&Markevitch,2003).With the aim of describing an evolutionary sequence of cluster formation,we have selected a small sample of galaxy clusters showing morphological evidence of ongoing merger activity.The sample was selected,on the basis of previous X-ray observations,from nearby clusters which entered theﬁeld of view of XMM-Newton.Here we summarise results for A1750and A3921,presented in detail in Belsole et al.(2004,2005),and we describe preliminary,new results for A2065.2Observations and data analysisIn Table1we list the main physical characteristics of the three clusters,to-gether with basic observation information.The A1750and A3921observations were very little contaminated by soft protonﬂares(details of the data prepa-ration can be found in Belsole et al.2004,2005).Unfortunately,emission from soft protons dominates the whole observation of A2065.This required an ad-hoc treatment for these data,which allowed us to model the background. A2065will be re-observed with XMM-Newton,and thus here we discuss only preliminary results.The background estimates were obtained using a blank-sky observation con-sisting of several high-latitude pointings with sources removed(Lumb et al. (2002)),and source and background events were corrected for vignetting us-ing the eviweight task in the Science Analysis System(SAS),enabling us to use the on-axis response matrices and eﬀective area.Table1Journal of observations.z0.0860.0720.096RA(J2000)13h30m52s15h22m42.6s22h49m38.6sDec(J2000)−01o50′27′′+27o43′′21′′−64o23′15′′Obs time(ks)302330N H(1020cm−2)a 2.39 2.95 2.94Global T(keV) 3.87(2.84)b 5.4 5.0(a)(b)(c)Fig.1.(a)XMM-Newton EPIC0.3-7.0keV energy band count image of A1750;(b)MOS1+MOS20.1-1.4keV energy band count image of A2065,the band chosen to enhance cluster emission above the strong particle background;(c) XMM-Newton EPIC0.3-10keV energy band count image of A3921.These are raw, non-background subtracted images.to A1750were taken into account simultaneously.We found that A1750does not show signiﬁcant residuals above aβ-model,in the region between the two clusters,but there is some excess emission in their cores.On the other hand,A3921shows large residuals to the west(other than those(a)(b)Fig.2.Surface brightness maps of A1750(a)and A3921(b)in the soft energy band (0.3-2.0keV).Contours are the residuals above a2Dβ-model which was subtracted from the image.Theﬁrst contour is at1σ.in the centre of the main cluster),indicating that a signiﬁcant substructure is present and it is related to one or both the brightest galaxies in that area.The two galaxies are at7arcmin(RA=22h48m49s,Dec=−64◦23′10′′(J2000))to the north-west(BG2)and at8arcmin(RA=22h49m04s,Dec=−64◦20′35′′(J2000))to the WNW(BG3)from the central brightest galaxy BG1(RA= 22h49m58s;Dec=−64◦25′46′′(J2000);see Belsole et al.2005for details).3.2Temperature distributionWe computed temperature maps for A1750and A3921by applying the multi-scale spectro-imagery algorithm described in(Bourdin et al.,2004).EMOS cameras only were used for this analysis.For A2065,we obtained an hardness ratio(HR)map once the particle contamination was modelled and subtracted from the images.This gives us qualitative,but signiﬁcant results about tem-perature substructure in this cluster.The temperature maps are shown in Figure3.All three clusters show signiﬁcant temperature variations,even in the cases where surface brightness substructures were not detected(A1750).As is con-ﬁrmed by extraction of spectra from discrete regions(Belsole et al.,2004), the region between A1750N and A1750C has a temperature which is signif-icantly(30%)hotter than the global temperature of either the two clusters. Additional temperature gradients are observed in A1750C,indicating that this(a)(b)(c)Fig.3.Temperature maps.The wavelet algorithm described in Bourdin et al.(2004) was used for A1750(a)and A3921(c).In the case of A2065(b)the image is an hardness ratio map(see text for details).cluster was already perturbed in the past.The temperature enhancement,coincident with the isophotal compression, observed in the A2065HR map points to a bow shock.We also detect a cool core,coincident with the peak of X-ray emission.The orientation of the striking hot temperature bar in A3921represents our strongest proof against the interpretation that this is a merger before close core passage.The temperature here varies between7and8keV,up to60% higher than the main cluster temperature.4DiscussionFor each cluster we summarise the results obtained in the previous sections.A1750(1)From the morphological analysis we found:oﬀ-centre cores and twist ofthe isophotes,but lack of substructure in the region between the two clusters.(2)the temperature increases weakly(∼30%),but signiﬁcantly,in the regionbetween the two clusters(3)strong temperature variations are observed within the main cluster(A1750C)and we also measured a discontinuity in the gas density proﬁle of order 20%.(4)we found a high entropy in the core of the two clusters(see Belsole et al.2004)The combination of these results lead us to interpret A1750as an ongoing merger between two clusters of similar masses,which have just started to in-teract at a real distance comparable to their virial radii.These two units will be in the compact phase(core passage)within1Gyr.However the temperature and density variations within A1750C itself are not explicable by the merger event occurring with A1750N and are intrinsic to this cluster.The most likely interpretation is that this is the signature of(a)previous merger(s).A1750is a good observational test case to be compared with numerical simulations, since it shows all of the signatures expected from gravitationally dominated processes.However,the fact that eﬀects of a previous merger are observed strongly supports the necessity of taking into account time dependent quan-tities,such us the relaxation time,in using galaxy clusters for cosmological purposes.A2065The surface brightness shows an isophotal compression towards the south-east,at∼80arcsec from the X-ray peak.The feature is accompanied by a forward shock in the same axis(NW-SE);moreover a surviving cooling coreis detected.Despite the limited quality of the data,A2065shows clear signatures of be-ing an ongoing merger in the compact phase,when the detection of strong shocks is the most favourable.This is similar to what is observed in numer-ical simulations of head-on collisions of merging clusters.The fact that thecore of the main cluster is cool suggests that this is probably the remnantof a cooling core,and thus the colliding object was probably of smaller mass (e.g.G´o mez et al.,2002).The new XMM-Newton observations of this object should allow us to put better constraint on these preliminary results.A3921Observational merger evidence includes:(1)Two peaks in both the X-ray emission and the galaxy distribution(Ferrari et al.,2005);(2)The hot bar is oriented parallel to the line joining the subclusters,it isnot orthogonal,as in the case of A1750;(3)The central regions of the main cluster and the subcluster to the westare strongly perturbed;(4)There is an oﬀ-set between the galaxy distribution of the smaller sub-cluster and the secondary X-ray peak(Belsole et al.,2005;Ferrari et al.,2005).This evidence cannot be interpreted as the result of a pre-merger.We are seeing maybe the best example of an X-ray observed oﬀ-axis post-merger.The subcluster has come from somewhere in the SE and is currently exiting towards the NW.The two merging units have diﬀerent masses,on a ratio of1to3(or1to5;see also Ferrari et al.2005).Comparison of these high qualityX-ray results with optical observations and numerical simulations yield an estimate of the age of the merger of order0.5Gyr after core passage.Oﬀ-axis mergers are more likely to occur than head-on ones,and they are more eﬃcientin mixing the gas via turbulence.If anything,A3921needs even deeper study,and the combination of these XMM-Newton data with upcoming Chandra observations will give further elements to our interpretation.5ConclusionsThe XMM-Newton observations of this small sample of galaxy clusters has conﬁrmed that the comparison of X-ray morphology and temperature is an excellent tool to understand the dynamical status of these objects.In the case of A1750and A3921,spectroscopy of discrete regions((Belsole et al., 2004,2005)has conﬁrmed the signiﬁcance of the temperature structure found with the multi-scale wavelet approach.The signiﬁcance of the temperature structure in A2065will be investigated with the upcoming XMM-Newton ob-servation.For the better studied cases we have conﬁrmed the previous inter-pretation of a recent merger(A1750),but have added new evidence suggesting that at least one of the subclusters is itself a merger remnant.Good quality X-ray data have allowed us to completely revolutionise the interpretation of the dynamical state of A3921.We can organise these clusters along an evolutionary path,where A1750rep-resents the beginning of a merging event,which in a timescale of order1.5 Gyr will be in the same state as A2065,when a bow shock is departing in the direction of motion after the two cores have collided,and the two(or maybe more)colliding objects are not physically separable anymore.Finally,A3921 represents the epoch when the secondary object has already passed the close core passage phase,it is exiting on the far side with respect to the direction of motion,and it will be accreted by the main cluster on a larger timescale (of order3-5Gyr).From so small a sample we cannot extract global conclusions on the popu-lation of merging galaxy clusters.However,this work supports the necessity of a wider investigation of the eﬀect of the physics of mergers on the global characteristics of galaxy clusters,especially if these objects are to be used to derive cosmological parameters.ReferencesBeers,T.C.,Gebhardt,K.,Forman,W.,Huchra,J.P.,Jones,C.,1991,AJ, 102,1581-1609,A dynamical analysis of twelve clusters of galaxies Belsole,E.,Pratt,G.W.,Sauvageot,J-L.,Bourdin,H.,2004,A&A,415,821-838,An XMM-Newton observation of the dynamically active binary cluster A1750Belsole,E.,Sauvageot,J-L.,Pratt,G.W.,Bourdin,H.,2005,A&A,in press, preprint astro-ph/0409544,An XMM-Newton observation of A3921:an oﬀ-axis mergerBourdin,H.,Sauvageot,J.-L.,Slezak,E.,Bijaoui,A.,Teyssier,R.,2004,A&A, 414,429-443,Temperature map computation for X-ray clusters of galaxiesCavaliere,A.,&Fusco-Femiano,R.,1976,A&A,49,137-144,X-rays from hot plasma in clusters of galaxiesDickey,J.M.,&Lockman F.J.,1990,ARA&A,28,215-261,HI in the Galaxy Ferrari,C.,Benoist,C.,Maurogordato,S.,Cappi,A.,Slezak,E.,2005,A&A, accepted,preprint:astro-ph/0409072,Dynamical state and star formation properties of the merging galaxy cluster Abell3921G´o mez,P.L.,Loken,C.,Roettiger,K.,Burns,J.O.,2002,ApJ,569,122-133, Do Cooling Flows Survive Cluster Mergers?Henry,J.P.,Finoguenov,A.,Briel,U.,2004,ApJ,615,181-195,Wideﬁeld X-ray temperature,pressure and entropy maps of A754Krivonos,R.A.,Vikhlinin,A.A.,Markevitch,M.L.,Pavlinsky,M.N.,2003, AstL,29,425-428,A Possible Shock Wave in the Intergalactic Medium of the Cluster of Galaxies A754Lumb,D.H.,Warwick,R.S.,Page,M.,De Luca,A.,2002,A&A,389,93-105, X-ray background measurements with XMM-Newton EPICMarkevitch et al.,2000,ApJ,541,542-549,Chandra Observation of Abell 2142:Survival of Dense Subcluster Cores in a MergerNeumann,D.M.,Lumb,D.H.,Pratt,G.W.,Briel,U.,2003,A&A,400,811-821,The dynamical state of the Coma cluster sith XMM-Newton Rowley,D.R.,Thomas,P.A.,Kay,S.T.,2004,MNRAS,352,508-522,The merger history of clusters and its eﬀect on the X-ray properties of the intr-acluster medium,Vikhlinin,A.A.,Markevitch,M.L.,2003,AstL,28,495-508,A Cold Front in the Galaxy Cluster A3667:Hydrodynamics,Heat Conduction and Magnetic Field in the Intergalactic Medium。

关于空间科学发展的一些思考

* 通信作者修改稿收到日期：2022年7月22日① 詹姆斯·韦布空间望远镜（James Webb Space Telescope ，JWST ）由美国、欧洲和加拿大联合研制，主要开展红外波段的宇宙起源、星系演化等观测研究，目前已公布了首批科学成果。

② 哈勃空间望远镜（Hubble Space Telescope ，HST ）于1990年4月发射，主要开展光学波段（延伸至近红外和近紫外）的宇宙学、星系、黑洞及恒星等观测研究。

③ 斯皮策太空望远镜（Spitzer Space Telescope ，SST ）于2003年8月发射，主要开展红外波段的宇宙学观测研究。

④ 钱德拉X 射线天文台（Chandra X-ray Observatory ，CXO ）于1999年7月发射，主要开展X 射线天文学研究。

⑤ 康普顿伽马射线天文台（Compton Gamma Ray Observatory ，CGRO ）于1991年4月发射，主要开展伽马射线天文学研究。

专题：中国空间科学――战略与突破Space Science in China: Strategy and Breakthrough 引用格式：顾逸东. 关于空间科学发展的一些思考. 中国科学院院刊, 2022, 37(8): 1031-1049.Gu Y D. Thoughts on space science development. Bulletin of Chinese Academy of Sciences, 2022, 37(8): 1031-1049. (in Chinese)关于空间科学发展的一些思考顾逸东*中国科学院空间应用工程与技术中心北京 100094摘要空间科学是利用空间飞行器探索宇宙和自然规律，开展特殊实验的大规模科学活动，为当代科学发展作出了重大贡献。

文章概述了国际空间科学的重大突破，分析了空间科学、载人探索活动的发展脉络和新趋势，回顾了我国空间科学近 20 年来的发展成就，总结了存在的问题和差距，强调空间科学的战略地位及其对我国科技、航天和国家长远发展的重要性。

人类首次探测到三颗类太阳恒星的恒星风

An international research team led by a researcher from the University of Vienna has for the first time directly detected stellar winds from three Sun-like stars by recording the X-ray emission from their astrospheres(星状体), and placed restrictions on the mass loss rate of the stars via their stellar winds.The researchers observed the spectral fingerprintsof the oxygen ions(离子) with XMM-Newton and were able to determine the quantity of oxygen and ultimately the total mass of stellar wind emitted by the stars. For the three stars with detected astrospheres, named 70 Ophiuchi, epsilon Eridani, and 61 Cygni, the researchers estimated their mass loss rates to be 66.5±11.1, 15.6±4.4, and 9.6±4.1 times the solar mass loss rate, respectively.人类首次探测到三颗类太阳恒星的恒星风●南昌民德学校　谢钰洪1. What’s worrying us when it comes to wireless devices?A. Shockingly wide applications.B. Unsustainable power support.C. Poor practical adaptability.D. Quick power consumption.2. What’s the main idea of Paragraph 3?A. The working principle of the PEC.B. The composition of the PEC.C. The performance of the new material.D. The complex design of the device.3. What’s the function of the electric double layer?A. To stimulate ions.B. To produce energy.C. To preserve electricity.D. To adjust the system.4. What can be inferred from what Roundy said?A. All your devices should be closely watched.B. Smart watcher’s sensors are better data senders.C. The new device is only workable to some sensors.D. The new device’s sensors can transmit big data.参考答案1. B 。

窄线赛弗特I型星系中可能具有吸积盘热辐射起源的软X射线超

窄线赛弗特I型星系中可能具有吸积盘热辐射起源的软X射线超李晔【摘要】理论上,低质量高吸积率活动星系核的标准吸积盘黑体热辐射有可能形成软X射线波段的超出.然而,活动星系核的软X射线超温度为0.1 keV左右,显著高于标准吸积盘理论预言的最高有效温度.只有Yuan等人在2010年报道的RXJ1633+4718是一个例外,其软超温度为32.5±-8.06.0eV,显著小于0.1 keV,并与吸积盘理论预言的温度相符；此外,其光度也符合吸积盘的理论预期.因此该软X射线超很可能起源于吸积盘的热辐射.这一类源对于吸积盘理论和软X射线超的研究都有重要意义.文中使用ROSAT PSPC数据,在所有已知的窄线赛弗特I型星系(Narrow Line Seyfert I galaxies,NLSIs)中寻找类似于RXJ1633+4718的源,即软X射线超可能起源于吸积盘热辐射的源.分析了150个源的245条光谱,其中58个源的90条光谱具有显著的软X射线超.样本中软X射线超温度分布的平均值-Tsoft为0.10 keV,标准偏差为0.03 keV.除了RX J1633+4718外,只有3个源的软X射线超温度小于60 eV,而且温度接近吸积盘最高温度.这个结果表明,类似于RX J1633+4718的具有低温软X射线超的源很稀少.最后,在具有黑洞质量估计的26个源中,软超温度与黑洞质量MBH、爱丁顿比L/LEdd及预期吸积盘最高温度Tmax都没有明显的关联性,与前人的结果一致.%It is theoretically expected that the thermal radiation from accretion disks of Active Galactic Nuclei (AGNs) with low black hole mass would contribute to the soft X-ray band and result in an excess emission. However, the temperatures of the observed soft X-ray excess of AGNs are typically around 0.1 keV, significantly higher than those predicted by the standard disk model. There is only oneexception found so far, R.X J1633+4718, which has a significantly low soft excess temperature of 32.5l+8.0-6.0 eV, consistent with the maximum temperature as well as the luminosity of the predicted disk emission(Yuan et al. 2010).rnIn this paper, we search for RX J1633 like AGNs, whose soft X-ray excess temperature is close to the maximum temperature of accretion disk, in all known Narrow Line Seyfert I galaxies (NLSIs) observed with ROSAT PSPC in the archive. We analyzed the PSPC spectra of 150 NLSIs (245 observations) in total, 58 (90 observations) of which show significantly soft X-ray excess emission. In addition to RX J1633+4718, only 3 objects (4 observations) show soft X-ray excess temperatures less than 60 eV and close to the theoretical prediction of the disk emission. They form a sample of candidates for future follow-up studies. Our result indicates that RX J1633+4718 like objects are very rare. The correlations between the temperature of the soft X-ray excess and black hole mass as well as the Eddington ratio are also investigated, and no correlation is found. This confirms the previous results, and implies that the soft X-ray excess of most AGNs should originate from other mechanisms rather than the thermal emission of the accretion disk.【期刊名称】《天文学进展》【年(卷),期】2013(031)001【总页数】14页(P110-123)【关键词】活动星系核;吸积盘X射线热辐射;软X射线超【作者】李晔【作者单位】中国科学院国家天文台/云南天文台,昆明650011;中国科学院大学,北京100049;中国科学院国家天文台,北京100012【正文语种】中文【中图分类】P157.6活动星系核中心存在超大质量黑洞，物质落入中心黑洞的过程中形成吸积流，并产生辐射[1-3]。

活动星系核X射线本征谱指数与爱丁顿比关系的研究进展

活动星系核X射线本征谱指数与爱丁顿比关系的研究进展倪嘉阳;薛永泉【摘要】活动星系核是中央核区有剧烈活动的(河外)星系总称.随着观测技术不断进步,人们对活动星系核的研究越来越多,对其理解也越来越深刻.总结整理了近年来对活动星系核X射线本征谱指数与爱丁顿比关系的观测结果,揭示出如下V形关系图像:随着爱丁顿比由大变小,X射线本征谱指数与爱丁顿比由存在正相关关系,转变为存在负相关关系.一般认为,这一观测现象反映了随着吸积率的降低,黑洞吸积模式发生了变化,由高吸积率时的标准薄盘吸积变为低吸积率时的辐射无效吸积流.这表明,基于标准薄盘的最基本的活动星系核统一模型虽然能够成功地解释较高光度活动星系核的很多观测现象,但却需要做一定的修正,以解释低光度活动星系核的一些观测性质.同时,将来有希望利用这一相关关系估算活动星系核一些重要参数,如中央超大黑洞质量、吸积率等,从而帮助人们更好地理解活动星系核的辐射机制和演化过程.最后对这一领域的研究进行了总结与展望.%Active galactic nuclei (AGN) are galaxies that have strong activities at their centers. Due to continual progress in observational technologies, a wide variety ofresearches on AGN have been carried out, leading to profound understanding of the nature of AGN. This paper collects and summarizes recent observational results about the correlation between X-ray intrinsic photon index and Eddington ratio in AGN, which present a V-shape correlation:as the Eddington ratio decreases continuously, such a correlation changes from being positive into being negative. It is generally believed that these observational results reflect that, as the accretion rate decreases, the black hole accretion mode changes as well, from theoriginal standard thin disk into radiatively ine?cient accretion flow. This indicates that the very basic AGN unification models based on the standard thin accretion disk have to be modified in order to well explain some observed properties of the low-luminosity AGN, although those models enjoy great success in explaining many observations of the AGN with higher luminosities. In the future, it is possible to use this correlation (if refined further) to estimate some important AGN parameters such as the black hole mass and accretion rate, which will facilitate our better understanding of the radiation mechanisms and evolution of AGN. Finally, this paper is concluded with some future prospects in this research field.【期刊名称】《天文学进展》【年(卷),期】2017(035)004【总页数】19页(P398-416)【关键词】活动星系核;X射线天文学;吸积物理【作者】倪嘉阳;薛永泉【作者单位】中国科学技术大学天文学系,合肥 230026;中国科学院星系与宇宙学重点实验室,合肥 230026;中国科学技术大学天文学系,合肥 230026;中国科学院星系与宇宙学重点实验室,合肥 230026【正文语种】中文【中图分类】P157.6活动星系核(active galactic nucleus，简称AGN)是宇宙中光度最大的一类天体，其在整个电磁波段都有很强的辐射。

首都师范大学入选“国家哲学社会科学成果文库”简介：《美国西部牛仔研究》

学、艺术、影视作品和各种商业广告中，牛仔被神化、偶像化，致使牛仔的真实历史严被遮
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美国西部牛仔研究》系在国家社科基金项目优秀结项成果的基础上修改戍书，２１）１书中部分内容以沧丈形式发表在
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成为一种国际现象但流行文化中的美国西部牛
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美国西；牛仔在美国西部史特别是西部牧
业开发历史上具有极其重要的地位．牛仔被美国
人视为开发西部的英雄，是征服西部荒野、开拓边疆的先锋，是美国精神的象征现今，牛仔的
／光明日报）史学月刊等报刊上，或在重要学术会议上宣读，产生了广泛影响程西；－＇３先生
评论说： “本书可谓全面、系统研究美国西部牛仔的一部拓荒之作，填补了美国西部开发史研究
的一个啦要空白 ”
“美国的非裔牛仔 ” 集中论述了美国西部牧区的
非裔牛仔历史概况第三章 “美国的牛仔罢工” 论析了美国历史上爆发的两次牛仔罢工及其意
０＾ ‘ 屑
义第四章 “牛仔的历史评价与热点论题 ” 主要概括和分析了牛仔的历史地位和贡献，并对当今牛仔学术研究中的几个热点论题予以评介和论析作者认为：美国西部牛仔是牧场主和牧牛公司

马来西亚华文学校华语管理与华人认同

马来西亚华文学校华语管理与华人认同*王晓梅（厦门大学马来西亚分校　中文系　马来西亚　雪邦　43900）提要马来西亚华人极力通过华文学校（华文独立中学与国民型华文中学）来维护他们的语言与文化，表现出强烈的华人认同。

本文采用认同管理模型与语言管理模型，以槟城州4所华文中学为例，讨论马来西亚的华人认同与华文学校华语管理之间的关系，分析这些华校如何通过语言管理（主要是华语）来强化华人认同。

数据采集主要通过面对面的访谈，涉及的访谈对象包括学校的管理层、教师与学生。

研究结果显示两类学校的华人认同与针对华语的语言规划措施密切相关，而且教师与学生对于华语的重要性以及华语对华人认同的重要性的看法基本一致。

此外，华文学校通过华语而实现的华人认同的管理与马来西亚华人社会的认同规范也是一致的。

关键词华人认同；华文学校；语言管理；华语；马来西亚中图分类号H002　文献标识码A　文章编号2096-1014（2018）03-0058-09DOI10.19689/10-1361/h.20180306Huayu, Chinese Education and Identity: The Malaysian ExperienceWang XiaomeiAbstract Chinese Malaysians have thrived to maintain their language and culture through Chinese schools including both In-dependent Chinese Secondary Schools and National-type Chinese Secondary Schools, which is an indication of strong Chinese identity. This paper adopts the models of identity management and language management to discuss the relationship between Chinese identity, the management of Mandarin, and Chinese education, taking four Chinese schools in Penang as examples. The main purpose is to analyze how these Chinese schools strengthen students’ Chinese identity through language management (mainly Mandarin). The results of interviews administered among school authorities, teachers and students show that the Chinese identity in both types of schools is closely associated with language planning towards Mandarin. Moreover, teachers and students hold similar views towards the importance of Mandarin and its signiﬁ cance to Chinese identity. Besides, the management of Chinese identity through Mandarin by Chinese schools complies with the norm of identity in Malaysian Chinese community.Key words Chinese identity; Chinese schools; language management; Mandarin; Malaysia一、前　言马来西亚语境下语言、教育与认同的关系一直受到学者们的关注，他们从教育的角度（Lopez 作者简介：王晓梅，女，厦门大学马来西亚分校副教授，主要研究方向为社会语言学与全球华语。

XMM 牛顿卫星

·服务舱（SVM），其中装有卫星的各个分系统和为卫星提供必备资源的相关单元。此外，在SVM上还安装有 2块太阳阵翼、1个望远镜遮阳罩（TSS）和2个安装在伸缩杆上的S波段天线。
XMM-牛顿卫星的实际结构
搭载仪器
XMM-牛顿卫星重3.8吨，长10米，太阳能电池帆板展开后宽16米。卫星上搭载的主要科学仪器有：
简介
XMM-牛顿卫星是欧洲航天局1999年发射的一颗X射线天文卫星，具有极高的谱分辨本领。XMM-牛顿计划始于 1984年，1997年3月开始建造，原名为“高通量X射线分光任务”（The High-Throughput X-ray Spectroscopy Mission），为纪念英国著名物理学家牛顿而命名XMM-牛顿卫星，其中XMM是X-ray MultiMirror Mission（X射线多镜面任务）的缩写。1999年12月10日，XMM-牛顿卫星在法属圭亚那的库鲁发射场用欧洲功率最大的火箭阿里安-5型火箭发射升空。它的轨道是椭圆形，近地点7,000公里，远地点114,000公里，轨道倾角40度，周期48小时。XMM-牛顿卫星原计划寿命为两年，但是将延期至最早2010年。
·3台掠射式X射线望远镜，每个望远镜由58个套筒组成，采用沃尔特Ⅰ型构造，最大的一层直径为70厘米，焦距为7.5米，总接收面积达到4,300平方厘米，由意大利科学家制造。
·3台欧洲光子成像照相机EPIC，工作能段为0.2-12keV，位于X射线望远镜的焦平面上，用于X射线成像、X 射线测光和中等分辨率分光。其中两台为MOS（金属氧化半导体）照相机，由7块不共面的CCD芯片组成，另一台为PN照相机，由12块共面的CCD芯片组成。
结构模块
XMM-牛顿卫星的设计紧紧围绕着X射线望远镜系统展开。它采用模块化结构，由四大部件组成：

2024年高校教师资格证之高等教育学自测提分题库加精品答案

2024年高校教师资格证之高等教育学自测提分题库加精品答案单选题（共200题）1、高等教育系统的主要层次是（）。

A.初级学院B.大学和专门学院C.研究生院D.函授学院【答案】 B2、“善于提出问题，启发学生积极思考”描述的是哪项大学教育原则（）。

A.科学性与思想性相统一原则B.传授知识与培养能力相统一原则C.教师主导作用与学生主动性相结合原则D.理论与实践相结合原则【答案】 C3、在教学过程中，教师既要对学生传授知识，又要发展学生的智力，培养学生的能力。

它体现了（）。

A.科学性和思想性相统一的原则B.面向全体与因材施教相结合的原则C.理论联系实际的原则D.传授知识与培养能力相统一的原则【答案】 D4、“自强不息，厚德载物”是著名高校（）的校训。

A.北京大学B.清华大学C.南开大学D.浙江大学【答案】 B5、根据“泰勒原理”，课程编制过程可以分为四个步骤或阶段，即确立目标，选择经验，组织经验和（）。

A.评价结果B.评价经验C.编辑教材D.整合经验【答案】 A6、下列选项中，为教育事业确立总的工作方向和奋斗目标的是（）。

A.教育目的B.教育方针C.人才培养方案D.教育政策【答案】 B7、要善于鼓励学生独立完成毕业设计或论文，这体现了高校教师在指导学生进行毕业设计或论文时的（）。

A.启发性原则B.教育性原则C.科学性原则D.因材施教原则【答案】 B8、教育与政治关系的核心是（）。

A.政治教育B.培养人才C.完善政体D.民主问题【答案】 D9、教师的根本职责是（）。

A.热爱学生B.严谨治学C.为人师表D.教书育人【答案】 D10、教师专业发展的关键是（）。

A.创新实践B.教研结合C.终身学习D.反思经验【答案】 D11、我国现行高校领导机制是（）。

A.校长分工负责制B.校务委员会负责制C.党委领导下的校长负责制D.教授委员会负责制【答案】 C12、教师管理的根本问题是（）。

A.激励教师勤奋工作B.营造和睦相处的氛围C.提升教师学历D.提高教师待遇【答案】 A13、不属于个案研究法的特征的是（）。

《科学美国人》与英特尔公司创建人之一戈登·穆尔访谈录(2)

《科学美国人》与英特尔公司创建人之一戈登穆尔访谈录(2)程永来
【期刊名称】《中国科教创新导刊》
【年(卷),期】1998(000)002
【摘要】《科学美国人》与英特尔公司创建人之一戈登·穆尔访谈录（２）程永来译方韧校戈登·穆尔，英特尔公司创始人之一和英特尔公司名誉主席，回顾了给他带来盛誉的穆尔定律。

从２５年前参与开发第一台“单芯片计算机”，到共同创建英特尔公司并使英特尔公司发展成一个半导体巨...
【总页数】2页(P28-29)
【作者】程永来
【作者单位】
【正文语种】中文
【中图分类】K837.126.1
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