Elemental abundances of Galactic bulge planetary nebulae from optical recombination lines

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。

小学下册英语第四单元真题(含答案)考试时间：90分钟（总分:120）A卷一、综合题(共计100题共100分)1. 听力题：They are _____ (cooking/eating) dinner together.2. 听力题：Heat is a form of ______ energy.3. 填空题：The rabbit has a cute ______ (鼻子).4. 填空题：The ______ (小鸟) sits on a branch.5. 听力题：I have a _____ of chocolate in my bag. (bar)6. 选择题：What do we call the season when flowers bloom?A. AutumnB. WinterC. SpringD. Summer7. 选择题：What is the capital of Portugal?A. LisbonB. MadridC. RomeD. Athens答案: A. Lisbon8. 选择题：What do you call a group of fish swimming together?A. SchoolB. PackC. FlockD. Herd答案: A9. 填空题：I enjoy visiting the ________ (历史博物馆).10. 填空题：The ant builds complex ______ (巢穴) underground.11. 填空题：A ____(green marketing) emphasizes sustainability in advertising.12. 选择题：What is the name of the princess in "Aladdin"?A. JasmineB. ArielC. BelleD. Cinderella答案:A13. 选择题：What do we call the time when we eat lunch?A. BreakfastB. LunchC. DinnerD. Snack答案: B14. 听力题：A _______ is the state of matter that has a fixed volume but not a fixed shape.15. 听力题：The symbol for thorium is _____.16. 听力题：We collaborate on ________ (tasks) regularly.17. 选择题：What is the primary function of the heart?A. To digest foodB. To pump bloodC. To breatheD. To think答案: B18. 选择题：What is the capital of Venezuela?A. CaracasB. BogotaC. QuitoD. Lima19. 选择题：What is the primary color of a stop sign?a. Blueb. Yellowc. Redd. Green答案:C20. 选择题：What is the capital of Lesotho?a. Maserub. Teyateyanengc. Maputsoed. Mohale's Hoek答案:a21. 填空题：The ________ (植物适应性) to climate is amazing.22. 填空题：Elements on the left side of the periodic table are usually __________ (金属).23. 听力题：I will _____ (help/helped) you with your homework.24. 听力题：The chemical formula for sodium phosphate is _____.25. 填空题：I can ______ (参与) in community service projects.26. 选择题：What do we call the time when leaves fall from trees?A. SpringB. SummerC. FallD. Winter答案: C27. grant application) requests funding for initiatives. 填空题：The ____28. 选择题：What is the term for a young sheep?A. CalfB. FoalC. LambD. Kid答案:C29. 填空题：Violets are _______ flowers that bloom in spring.30. 选择题：根据短文内容,选出Tom和 Chen Jie 书包里的物品。

五年级世界历史英语阅读理解25题1<背景文章>Ancient Egypt is one of the most fascinating civilizations in history. The Egyptians made many great achievements. One of their most famous creations is the pyramid. Pyramids were huge structures built as tombs for the pharaohs. They were made of large stones and took many years to build.Another amazing achievement of the Egyptians is the mummy. Mummies are the preserved bodies of dead people. The Egyptians believed that the body needed to be preserved so that the soul could return after death. They used special techniques to dry out the body and wrap it in linen bandages.The Egyptians also had a very advanced writing system called hieroglyphics. Hieroglyphics were pictures and symbols that represented words and ideas. Only a few people knew how to read and write hieroglyphics.Egypt was also known for its beautiful art and architecture. The temples and palaces were decorated with colorful paintings and carvings.1. The pyramids were built as ________ for the pharaohs.A. homesB. templesC. tombsD. palaces答案：C。

a r X i v :a s t r o -p h /0506238v 1 10 J u n 2005AJ,IN PRESSPreprint typeset using L A T E X style emulateapj v.11/26/04A COMPARISON OF ELEMENTAL ABUNDANCE RATIOS IN GLOBULAR CLUSTERS,FIELD STARS,AND DW ARFSPHEROIDAL GALAXIESB ARTON J.P RITZL AND K IM A.V ENNDepartment of Physics and Astronomy,Macalester College,1600Grand Ave.,Saint Paul,MN 55105email:pritzl@;venn@M IKE J.I RWINInstitute of Astronomy,University of Cambridge,Madingley Road,Cambridge,CB30HA,UKemail:mike@AJ,in pressABSTRACTWe have compiled a sample of globular clusters with high quality stellar abundances from the literature to compare to the chemistries of stars in the Galaxy and those in dwarf spheroidal galaxies.Of the 45globular clusters examined,29also have kinematic information.Most of the globular clusters belong to the Galactic halo,however a signficant number have disk kinematics or belong to the bulge.Focusing on the [α/Fe]and light r-process element ratios,we find that most globular cluster stars mimic those of the field stars of similar metallicities,and neither clearly resembles the presently available stellar abundances in the dwarf galaxies (including the globular clusters in the Large Magellanic Cloud).The exceptions to these general elemental ratio comparisons are already known in the literature,e.g.,ωCentauri,Palomar 12,and Terzan 7associated with the Sagittarius remnant,and Ruprecht 106which has a high radial velocity and low [α/Fe]ratio.A few other globular clusters show more marginal peculiarities.The most notable one being the halo cluster M68which has a high Galactocentric rotational velocity,a slightly younger age,and a unique [Si/Ti]ratio.The [Si/Ti]ratios decrease with increasing [Fe/H]at intermediate metallicities,which is consistent with very massive stars playing a larger role in the early chemical evolution of the Galaxy.The chemical similarities between globular clusters and field stars with [Fe/H]≤−1.0suggests a shared chemical history in a well mixed early Galaxy.The differences to the published chemistries of stars in the dwarf spheroidal galaxies suggests that neither the globular clusters,halo stars,nor thick disk stars had their origins in small isolated systems like the present-day Milky Way dwarf satellites.Subject headings:Globular clusters:general —stars:abundances1.INTRODUCTIONIf the Galaxy formed primarily through continuous merg-ing of small dwarf systems as demanded by cold dark mat-ter scenarios (e.g.,Navarro,Frenk,&White 1997;Klypin et al.1999;Moore et al.1999;Bullock,Kravstov,&Weinberg 2001;Boily et al.2004),possibly even until z =1,then we ought to be able to trace this assembly through the ratio of the elemental abundances in stars.This is because the color-magnitude diagrams of dwarf galaxies in the Local Group have shown a wide variety of star formation histories and physical properties (e.g.,Grebel 1997;Mateo 1998;Tolstoy et al.1998;Dolphin et al.2003,Skillman et al.2003),which are expected to lead to differences in chemical evolution,as predicted by many evolution scenarios (e.g.,Lanfranchi &Matteucci 2004;Pagel &Tautvai˘s ien˙e 1998).A variety of element ratios are expected in the stars of different ages and dwarf galaxy origin.Therefore,“chemical-tagging”(a term from Freeman &Bland-Hawthorn,2002)of stellar popula-tions in the Galaxy should be possible if significant formation occurred through hierarchical merging.The detailed and recent stellar abundances in the small-est,dwarf spheroidal galaxies (dSphs;Shetrone et al.2001,2003;Tolstoy et al.2003;Geisler et al.2005),have shown their stars tend to have lower [α/Fe]abundance ratios than similar metallicity Galactic field stars.The [α/Fe]abundance ratio is an important tracer of the relative contributions of SNe II/Ia since only SNe II contribute to the production of the αelements,whereas both contribute to iron.The [α/Fe]ratio is particularly sensitive to this difference at higher metallicities,when SNe Ia have begun to make a significant contribution to the chemical evolution of the system;the Sgr dwarf galaxy remnant is an excellent example of this,where the [α/Fe]trend is similar to that of Galactic halo stars,however offset such that it appears that SNe Ia have contributed significant amounts of iron at lower metallicities and [α/Fe]is signifi-cantly subsolar by [Fe/H]=0(Bonifacio et al.2004).However some galaxies,e.g.,Draco,show low [α/Fe]ratios even at the lowest metallicities ([Fe/H]∼−3.0;Shetrone,Côté,&Sargent 2001),which may suggest that metal-dependent out-flows of SNe II ejecta are also important.Many metal-poor dSph stars also show low [Y/Eu]abun-dance ratios,which depends on the r-and s-process contri-butions from SNe II and AGB stars.Because of the slower chemical evolution in the small dwarf galaxies,s-process en-richments from metal-poor AGB stars are important and these stars have lower yields of light s-process elements,such as Y (Travaglio et al.2004).In dSph stars,the low Y is not mim-iced by low Ba (as predicted),which results in high [Ba/Y]ratios compared to the Galactic field stars (Venn et al.2004).This is most evident around metallicities of [Fe/H]∼−2.0,be-fore the more metal-rich AGB stars can also contribute to the s-process yields,even dominate the total contribution which can bring the [Ba/Y]ratio,and others,into agreement with the Galactic comparison stars (e.g.,one metal-rich RGB/AGB star in Fornax has lower [Ba/Y]than the metal-poor RGB stars2Pritzl,Venn,&Irwinin Fornax,in excellent agreement with Galactic comparison field stars).The chemical comparisons of a large sample of Galactic field stars were compared to those in the small dSph galax-ies by Venn et al.(2004).There the stars were divided into Galactic components based on their kinematics.While there is some marginal evidence that Galactic halo stars with ex-treme retrograde orbits overlap in[α/Fe]with the stars in dSphs,this was not evident in the[Ba/Y]ratios.Thus,it ap-pears that no significant component of the Galactic halo,nor the Galactic thick disk which was also examined,could have formed from the mergers of these small dwarf galaxies.How-ever,these comparisons were based on the available detailed abundances in seven dSph galaxies,and primarily from stars in their centralfields.Early results from the much broader CaT survey in three dSphs(Sculptor,Fornax,and Sextans from the DART survey;e.g.,see Tolstoy et al.2004for Sculp-tor)shows that the stars in their centralfields tend to have higher metallicities than most of the stars out near their tidal radii.It is possible then that these outer dSphfield stars have different chemical abundance ratios,which may be more sim-ilar to the metal-poor stars in the Galactic halo;detailed abun-dance ratios for these outer metal-poor dSph stars are pend-ing.Aside from the selection of stars in the dSphs,the com-parisons examined by Venn et al.(2004)also did not rule out early merging,before the dwarf galaxies had a chance to have a significant and unique chemical evolution history;merging with larger dwarf galaxies was also not ruled out.It is also interesting to note that another important physical difference between the large and small dSphs galaxies is that the larger dwarf galaxies have higher masses and contain globular clus-ters(GCs;van den Bergh2000).In this paper,we ask how do the GCsfit into this scheme? Can the GCs be used as a test of the merging history of large dwarf galaxies in the formation of our Galaxy?If some of the Galactic GCs have extragalactic origins,then we may assume that they trace the merging history of the large dwarf galax-ies,given that GCs are only found in the larger Milky Way dwarf galaxies of the Magellanic Clouds,Sagittarius(Sgr), Fornax,and Canis Major(CMa).We caution that the op-posite is not true;not all massive dwarf galaxies have GCs. Therefore we can only say that if wefind an extragalactic GC that it must have come from a more massive dwarf galaxy. Presumably,differences in star formation histories,chemical evolution,and/or the initial conditions(mass,pre-enrichment, or environmental factors)between the various dwarf galax-ies could also leave discernable chemical imprints on their GCs that might be chemically tagged when merged into our Galaxy.It has been shown that the Sgr GCs Terzan7(Ter7) and Palomar12(Pal12)follow the abundance trends of the Sgrfield stars(e.g.,Cohen2004;Tautvai˘s ienéet al.2004). In addition,the GCs andfield stars analyzed in the Large Magellanic Cloud(LMC;Hill et al.2000;Hill2004)follow the same abundance trends that have lower[α/Fe]ratios(O, Ca,and Ti;with the possible exception of Mg)than similar metallicityfield stars in the Galaxy,which is similar to the field stars in the small dwarf galaxies.Thus,the question is whether there is a population of GCs in the Galaxy that have low[α/Fe]ratios(and possibly other chemical signatures) that could be interpreted as these clusters having formed in dwarf galaxies and later were captured through merging. There have been many papers that present and/or review the abundances in GCs,or compare a subset of GCs to one another and/or to Galacticfield stars(most recently Sneden,Ivans,&Fulbright2004;Gratton,Sneden,&Carretta2004). In general,the Galactic GCs show amazing uniformity in their [α/Fe]ratios(plateau levels∼+0.3dex),and they follow the abundance trends seen in the Galacticfield stars when plot-ted as a function of the[Fe/H]values.This seems incredi-ble in terms of understanding hierarchical galaxy formation. Carney(1996)showed that there is little or no relationship between[α/Fe]and age for the GCs.Given SNe Ia con-tributions lead to a decrease in the[α/Fe]ratios with in-creasing metallicity as seen in the Galacticfield stars,Car-ney argued that the lack of a similar turn-down for the GCs implies a lack of SNe Ia contributions to the GCs.There-fore,if the timescale of when the Type Ia supernovae sig-nificantly contribute to the interstellar medium is short,the “old”halo and disk GCs could not share a common chemi-cal history and one of the populations must have been later accreted.The exceptions include GCs associated with the Sgr dwarf galaxy,e.g.,Ter7and Pal12,which have lower [α/Fe]ratios than comparison Galactic stars(Cohen2004; Sbordone et al.2005;Tautvai˘s ien˙e et al.2004),and both have younger ages than typical Galatic halo GCs.Also,ωCen-tauri(Pancino et al.2002)and Ruprecht106(Rup106;Lin &Richer1992;Brown,Wallerstein,&Zucker1997)show peculiar chemical abundances and are thought to be captured clusters.But when it comes to GCs,the more common ques-tions have been related to the internal variations in CNO and NaMgAl observed in their red giant stars.These variations are attributed to a combination of initial composition differ-ences coupled with internal mixing mechanisms(Gratton et al.2004),though others have been exploring the possibilities and predictions of enrichments from early AGB stars during the cluster formation process(Cottrell&Da Costa1981).Be-cause these abundance anomalies are seen only in GCs and never infield stars,including in dwarf galaxyfield stars1, Shetrone et al.(1998,2001,2003)concluded that the GCs cannot have formed from the small dwarf galaxies,e.g.,dur-ing the merging event.However,GCs formed in large dwarf galaxies that later merged into the Galactic halo and survived has not been ruled out.In this paper,we re-examine the element ratios in as many Galactic GCs as available in the literature to compare with the Galacticfield stars and stars in the dwarf galaxies.The goal is to indentify candidate GCs that show signatures of extra-galactic origins either chemically or ideally both chemically and kinematically.While it has been found that smaller dwarf galaxies(such as the present-day Milky Way dSph satellites) have contributed little or nothing to the Galactic halo(at least from those stars as summarized in Venn et al.2004),the GCs should examine the contributions of larger dwarf galaxies. Thefield stars with detailed abundance ratios include only stars that are currently in the solar neighborhood,while the GCs sample all parts of the Galaxy,including the bulge.Of particular interest are the[α/Fe]ratios,as well as the light r-process ratios which have been successful in separating the Galacticfield stars from stars in the dSphs.We study the GC abundance ratios with regards to their kinematic populations having adopted the Galactocentric velocities(Π,Θ,W)from Dinescu et al.(1999a,b,2000,2001,2003,and private com-munications for updates)to separate the GCs into standard Galactic components(as discussed in§3).By noting which 1One report of a star in the Sgr dwarf galaxy remnant by Smecker-Hane& McWilliam(2005).However the metallicity and location of their metal-poor stars make them possible members of the M54GC(Bonifacio et al.2004).Elemental Abundance Comparisons3F IG .1.—Toomre diagram showing stellar and globular cluster populations using kinematic probabilities from velocity ellipsoids;thin disk (red),thick disk (green),halo (cyan),extreme retrograde stars (black).Field stars are points,while globular clusters are larger black circles filled with a color that refers to its phase space population.M54and Terzan 7are included in this plot with the velocity information for the Sagittarius dwarf galaxy remnant.Ruprecht 106is included after assuming a range in proper motions (see text).As expected,a little over half of the globular clusters have halo kinematics.See Venn et al.(2004)for field star data references.component of the Galaxy the clusters belong to,we also try to identify clusters that may have unusual abundance signatures for their kinematically assigned Galactic component.2.THE DATABASEDetailed elemental abundances are available for 45GCs;32clusters with ≥3stars analyzed.These clusters are listed in Table 1,along with the number of stars analysed in each cluster and references.The literature was searched for high-resolution abundances determined within the past 20years.The oldest analysis included here is by Gratton,Quarta,&Ortolani (1986)for three GCs.A majority of the analyses are from the past five years.We have found that studies from older than 1990used lower resolution spectra (∼15000).More modern studies that use lower resolution (18000or less)include McWilliam,Geisler,&Rich (1992),Mishen-ina,Panchuck,&Samus (2003),Cavallo,Suntzeff,&Pila-chowski (2004).We consider the abundance ratios coming from these studies as being less reliable.In most cases when abundance ratios based on higher resolution spectra are avail-able,those from lower resolution spectra are not included in the final mean ratios.Table 2lists the elemental abundances and their standard error of the mean for select elements from Mg to Eu for each of the GCs in Table 1.We adopted the solar abundances and log g f values used by the Lick-Texas group as the standard values (solar abundances,Grevesse &Sauval 1998;log g f ,see Table 3in Fulbright 2000for example).All abundance ratios were corrected to these values (observed -adopted)in an effort to standardize the different datasets 2.Details on the adjustments are given in §4.In Table 2,the final individual abundance ratios relative to Fe (shown in bold)were averaged together by weighted averages according to the number of2Unless no solar abundances and/or log g f values were given in the pub-lished article.stars that were used in each study.The [α/Fe]and comparison ratios,e.g.[Ba/Y],were calculated using the final weighted-mean ratios.The italicized ratios were not used in calculating the final weighted mean abundance ratios.The GCs in Table 2span a range of −2.38(M15)≤[Fe/H]≤−0.06(NGC 6528).In a majority of the cases,the [FeI/H]and [FeII/H]ratios were averaged when both were available.Similarly,the [Ti/Fe]ra-tio is the average of the [TiI/Fe]and [TiII/Fe]ratios when both were given for the cluster,otherwise it is the [TiI/Fe]ratio.Ideally,it would be best to match ionization states when deter-mining the Ti ratios,e.g.,[TiI/FeI]and [TiII/FeII].However,in most cases not enough information was presented to de-termine these values.Many papers give the final abundances [TiI/Fe]and [TiII/Fe]with no clear explanation how [FeI/H]and [FeII/H]were averaged together to determine [Fe/H].As a result,we used the mean [Fe/H]ratio when determining the Ti ratios.Our selection of elements was based on those deter-mined regularly in GC stars,and those that have been useful in the analysis of stars in dSph galaxies.Variations in CNO,Na,Mg,and Al are regularly found in clusters from star-to-star,and are usually attributed to mixing with CNO-cycled gas (e.g.,Kraft 1994;Sneden 2000;Free-man &Bland-Hawthorn 2002;Sneden et al.2004;Gratton et al.2004).We have avoided these elements in this paper (with the exception of magnesium 3given its relevance as a true αelement indicator)because we are primarily concerned with global signatures of galaxy formation and not stellar or chem-ical evolution within GCs.For abundances in Galactic field stars,we have adopted those in Table 2in Venn et al.(2004;see references therein).The abundances in seven dSph galaxies are from Shetrone et al.(2001,2003)and Geisler et al.(2005).Stellar abundances for 26stars in the Sgr dwarf galaxy remnant have also been3Where mixing with CNO-cycled gas typically results in ∆Mg ≤0.3dex (see discussion in §5.1)4Pritzl,Venn,&IrwinF IG .2.—Variations in [Fe/H]and [α/Fe]as a function of the Galactocentric rotational velocity (Θ)using the same symbols as in Fig.1.The field stars show a large range in the thick disk and halo components,which also significantly overlap one another.As in Fig.1,M54and Terzan 7are included in this plot with the velocity information for the Sagittarius dwarf galaxy remnant,and Ruprecht 106is included after assuming a range in proper motions (reflected by the error bar;see text).The globular clusters tend to have a more narrow range in [Fe/H]and [α/Fe](where [α/Fe]=[(Ca+Ti)/2Fe]),with only a few exceptions which are discussed in §3and §5.added from Bonifacio et al.(2000;2004)and Smecker-Hane &McWilliam (2005).3.GLOBULAR CLUSTER KINEMATIC ASSIGNMENTSWe have determined the Galactic stellar population com-ponent for 29GCs based on a calculation of their phase space.Phase space distribution functions have been deter-mined using the Galactocentric velocity vector components (V(Π,Θ,W )in km s −1)from Dinescu et al.(1999ab,2000,2001,2003,and private communications for updates)and the Galactocentric positional vector components (R(X,Y ,Z)in kpc;Harris 1996).Firstly,we computed the probability for a GC to be associated with the Galactic thin disk,thick disk,or halo from its velocity vector using a standard Bayesian classi-fication scheme and Galactic Gaussian velocity ellipsoid com-ponents from Dehnen &Binney (1998;thin disk),Soubiran,Bienaymé,&Siebert (2003;thick disk),and Chiba &Beers (2000;halo).This method is the same as was used for the Galactic field stars by Venn et al.(2004).Secondly,the prob-ability for a GCs to be associated with these Galactic compo-nents was determined from its positional vector.For this,a standard Galactic model (Robin,Reyleé,Derrière,&Picaud 2003)was chosen with a maximum halo extent of R =150kpc,a softening parameter ρ=1.0(which is used to stop the density of the halo from going to infinite at the Galactic Cen-ter),and adopting an asympotitic r −3halo profile.This was not necessary for the Galactic field stars studied by Venn et al.(2004)since it is a good assumption that those stars are in the solar neighbourhood quadrant.The kinematic and posi-tional probabilities were combined for a phase space deter-mination of the final probabilities for each GC to be associ-ated with each Galactic component.Finally,we allocated any GCs within R =2.7kpc of the Galactic Center to a bulge component unless the kinematics placed it in another compo-nent,and did not probe more deeply into variations between bulge GCs and/or those associated with a bar (e.g.,Dinescu et al.2003).Table 3lists the velocity and positional vector components for each GC,as well as which Galactic component they be-long to.We have also added kinematic information for M54and Ter 7which are embedded in the Sgr dwarf galaxy rem-nant;thus,on the assumption that these clusters are associ-ated with the Sgr remnant,then we have adopted the known kinematics for Sgr (Ibata et al.1997)as a reasonable approxi-mation to their Π,Θ,W velocities.This allows us to highlight these clusters in our abundance ratios and kinematic analyses.Rup 106is another cluster that does not have kinematic infor-mation,however it has a large galactocentric radial velocity (−232km s −1;Harris 1996).Adopting a plausible range of proper motions (0",±1"in RA and DEC),we investigate its potential space velocities (UVW and Π,Θ,W ).In all cases,the Galactocentric radial velocity,Π,is quite large (−200to −330km s −1;see Table 4)implying Rup 106is a member of the Galactic halo,possibly on a plunging orbit that would be consistent with a captured cluster.The final column in Table 3lists the GC classifications from (Mackey &Gilmore 2004)which are based on the physical properties of GCs.While overall there is good agreement between the two classifica-tions,there are some key differences especially for the clusters which have thick disk kinematics,while according to Mackey &Gilmore they have halo-like properties.A Toomre diagram for the GCs is shown in Figure 1,where the Galactocentric rotational velocity Θis plotted against T (T 2=Π2+W 2).To be consistent with the field stars exam-ined by Venn et al.(2004),we plot the GCs as colored points according to the assigned Galactic component (e.g.,cyan =halo,green =thick disk,red =thin disk).GCs assigned to the bulge are not shown in Fig.1.Only three clusters stand out in this diagram;Pal 12has thin disk kinematics however its positional vector places it in the halo (thus cyan colored GC amongst the thin disk field stars),M68which has an unusually high Galactocentric rotational velocity (Θ=+303km s −1),and M22is a thick disk cluster with an unusually high T compo-nent (+212km s −1).Figure 2shows the distribution in [Fe/H]and [α/Fe]vs.Θfor the GCs compared with the field stars in the solar neigh-borhood.Venn et al.(2004)commented on the large and over-lapping distribution in these abundances for the field stars in each of the Galactic components.For the [α/Fe]plot,we have averaged the mean abundances of Ca and Ti per GC (note that this differs slightly from the field star analysis by Venn et al.2004where [α/Fe]represented an average of Mg,Ca,and Ti).The most outstanding GCs are Pal 12and Rup 106,with very low [α/Fe]ratios.Of course,Pal 12is associatedElemental Abundance Comparisons5with the Sgr dwarf galaxy remnant and Rup106is thought to have been captured possibly from the Magellanic Clouds (Lin&Richer1992).For the thick disk clusters,it is inter-esting to see that they are clustered both toward lower[Fe/H] and higher[α/Fe].Although they are within the range of the thick diskfield stars,the thick disk GCs do not show as wide of a spread in metallicity.In any case,we conclude that the [α/Fe]ratios from the GCs in each kinematic component are in good to excellent agreement with those of thefield stars in the same Galactic component.4.ABUNDANCE RATIO CORRECTIONS FOR GLOBULARCLUSTERSMost of the data on GC stellar abundances comes from the Lick-Texas group.Therefore we have adopted the solar abun-dances they used(Grevesse&Sauval1998),along with their log g f values(see Table3of Fulbright2000for references), as thefiducial standards on which all abundances are adjusted. Differences in solar abundances and log g f are in the sense of the value from the reference paper minus our adopted values.A number of clusters were not corrected because either their solar abundances and log g f values matched the adopted ones (no adjustments were made for differences of less than about 0.04dex)or no values were given.These GCs were noted in Table1.In the following we note each cluster where adjust-ments were made to the abundance ratios.No adjustments for differences in hyperfine splitting corrections have been made because such abundance ratio adjustments are not as simple as those for solar abundances and log g f.NGC104(47Tucanae):Although the values from Grat-ton,Quarta,&Ortolani(1986)are not used in thefinal mean abundances,we adjust them to match our adopted log g f val-ues.There are no solar abundances given in their paper,so no adjustments were made.[Mg/Fe],[Si/Fe],[Ca/Fe],[TiI/Fe], and[Ba/Fe]are corrected by+0.22,+0.21,+0.20,+0.12,and −0.86,respectively,assuming the log g f values are the same as in Gratton(1987).The values from Brown&Wallerstein (1992)are not used in thefinal mean abundances,but we ad-justed them to match our adopted solar abundances and log g f values.Since the logǫvalues are given,the ratios are de-termined directly from those values using our adopted solar abundances.From the log g f values,[Mg/Fe],[Si/Fe],and [TiII/Fe]have been corrected by+0.09,+0.05,and+0.25,re-spectively.To avaoid confusion,in Table3we have combined the results from Gratton et al.(2001)and James et al.(2004a) since they are both part of a collaboration that studies the same stars,but different element ratios.The abundances from James et al.are weighted means of the turnoff and subgiant star ratios.NGC288,NGC362,NGC5897,NGC6352,and NGC6362:No solar abundance values were given in Gratton (1987),so the only adjustments that were made were due to the log g f values.[Mg/Fe],[Si/Fe],[Ca/Fe],[TiI/Fe],and [Ba/Fe]are corrected by+0.22,+0.21,+0.20,+0.12,and −0.86.For NGC288and NGC362,the results from Gratton (1987)are superceeded by those in Shetrone&Keane(2000). NGC2298:No changes are needed due to the solar abun-dance values in McWilliam,Geisler,&Rich(1992).Only [La/Fe]is adjusted by+0.09due to the log g f values.NGC3201:We do not use the Gratton&Ortolani(1989) ratios because two of the three stars are redone by Gonzalez& Wallerstein(1998).Examining the solar abundances,adjust-ments in[Mg/Fe],[TiI/Fe],and[Ba/Fe]needed to be made for the Gonzalez&Wallerstein abundance ratios.Further correc-tions need to be made due to differences in the log g f values for[Mg/Fe],[TiI/Fe],and[Eu/Fe].It happens that the adjust-ments for the differences in the solar abundances and log g f cancel each other out for[Mg/Fe]and[TiI/Fe]so that no ad-justments are necessary.The total changes for[Ba/Fe]and [Eu/Fe]are+0.23and+0.17,respectively.In determing the Gonzalez&Wallerstein ratios,we have chosen to use the high resolution abundances from1991and1994for those stars ob-served multiple times.NGC4590(M68):The results from Shetrone et al.(2003) and Lee,Carney,&Habgood(2004)superceed those from Gratton&Ortolani(1989).The abundance ratios derived from photometric values in Lee,Carney,&Habgood were used.NGC5273(M3):The results of Kraft et al.(1993,1995) were redone by Sneden et al.(2004).In addition,the three stars in Shetrone,Côté,&Sargent(2001)are also done in either Sneden et al.or Cohen&Melendez(2005a).Therefore, we did not use their abundance ratios in thefinal weighted-mean ratios.For the Cohen&Melendez study,the[Fe/H] values were adjusted by−0.07due to a difference in the solar abundance value.The abundance ratios[Mg/Fe],[Ca/Fe],and [Eu/Fe]were adjusted due to differences in the log g f values by−0.10,+0.16,and−0.08.Therefore thefinal abundance ratios for M3are a weighted mean of those from Sneden et al.(2004)and Cohen&Melendez(2005a).NGC5466:There are no changes to the NGC5466abun-dance ratios.We note that the only available chemical abun-dances come from the anomalous Cepheid in this cluster.This type of variable star can derive from either binary mass trans-fer or younger stars(Demarque&Hirshfeld1975;Norris& Zinn1975;Renzini,Mengel,&Sweigart1977).Because of the star’s variability,we caution about the reliability of the abundance ratios from this single star even though they are consistent with other stars and clusters of similar metallicity. NGC5904(M5):Although the values from Gratton, Quarta,&Ortolani(1986)are not used in thefinal mean abun-dances,we adjust them to match our adopted log g f values. There are no solar abundances given in the paper,so no ad-justments can be made.[Mg/Fe],[Si/Fe],[Ca/Fe],[TiI/Fe], and[Ba/Fe]are corrected by+0.22,+0.21,+0.20,+0.12,and −0.86,respectively,assuming the log g f values are the same as in Gratton(1987).The ratios from Sneden et al.(1992) are not used because the stars were reanalyzed by Ivans et al.(2001).The asymptotic giant branch stars in the Ivans et al.study were not included in the mean abundance ratios.For the Ramírez&Cohen(2003)abundance ratios,adjustments are necessary for both the solar abundances and the log g f values.For the solar abundances,[Mg/Fe],[Ca/Fe],[TiI/Fe], [FeI/H],[FeII/H],and[La/Fe]need to be adjusted by−0.12,−0.22,−0.07,−0.08,−0.05,and−0.08,respectively.For the log g f values,[Mg/Fe],[Ca/Fe],and[Eu/Fe]need corrections of−0.08,+0.16,and−0.08,respectively.This leads to a to-tal adjustments of−0.20,−0.06,−0.07,−0.08,−0.05,−0.08, and−0.08for[Mg/Fe],[Ca/Fe],[TiI/Fe],[FeI/H],[FeII/H], [La/Fe],and[Eu/Fe].After all of the correction were made, the Ivans et al.and Ramírez&Cohen ratios were averaged together by a weighted mean.NGC6093(M80):The solar abundances in Cavallo, Suntzeff,&Pilachowski(2004)matched our adopted val-ues.For the log g f values,correction were made to[Ca/Fe], [TiII/Fe],and[Eu/Fe]by+0.13,−0.06,and+0.22,respec-tively.。

剑桥商务英语听说星系The Milky Way GalaxyThe Milky Way is the galaxy that contains our Solar System, with the Earth and Sun. This galaxy is a vast, spinning collection of stars, planets, dust and gas, held together by gravity. It is just one of hundreds of billions of galaxies in the observable universe.The Milky Way galaxy is estimated to contain 100-400 billion stars and have a diameter between 100,000 and 180,000 light-years. It is the second-largest galaxy in the Local Group, with the Andromeda Galaxy being larger. As with other spiral galaxies, the Milky Way has a central bulge surrounded by a rotating disk of gas, dust and stars. This disk is approximately 13 billion years old and contains population I and population II stars.The solar system is located about 25,000 to 28,000 light-years from the galactic center, on the inner edge of one of the spiral-shaped concentrations of gas and dust called the Orion Arm. The stars in the Milky Way appear to form several distinct components including the bulge, the disk, and the halo. These components are made of different types of stars, and differ in their ages and their chemicalabundances.The Milky Way galaxy is part of the Local Group, a group of more than 50 galaxies, including the Andromeda Galaxy and several dwarf galaxies. The Local Group in turn is part of the Virgo Supercluster, a giant structure of thousands of galaxies. The Milky Way and Andromeda Galaxy are moving towards each other and are expected to collide in about 4.5 billion years, although the likelihood of any actual collisions between the stars themselves is negligible.The Milky Way has several major arms that spiral from the galactic bulge, as well as minor spurs. The best known are the Perseus Arm and the Sagittarius Arm. The Sun and its solar system are located between two of these spiral arms, known as the Local Bubble. There are believed to be four major spiral arms, as well as several smaller segments of spiral arms.The nature of the Milky Way's bar and spiral structure is still a matter of active research, with the latest research contradicting the previous theories. The Milky Way may have a prominent central bar structure, and its shape may be best described as a barred spiral galaxy. The disk of the Milky Way has a diameter of about 100,000 light-years. The galactic halo is a spherical component of the galaxy that extends outward from the galactic disk, as far as 200,000 light-years from the galactic center.The disk of the Milky Way Galaxy is marked by the presence of a supermassive black hole known as Sagittarius A*, which is located at the very center of the Galaxy. This black hole has a mass four million times greater than the mass of the Sun. The Milky Way's bar is thought to be about 27,000 light-years long and may be made up of older red stars.The Milky Way is moving with respect to the cosmic microwave background radiation in the direction of the constellation Hydra with a speed of 552 ± 6 km/s. The Milky Way is a spiral galaxy that has undergone major mergers with several smaller galaxies in its distant past. This is evidenced by studies of the stellar halo, which contains globular clusters and streams of stars that were torn from those smaller galaxies.The Milky Way is estimated to contain 100–400 billion stars. Most stars are within the disk and bulge, while the galactic halo is sparsely populated with stars and globular clusters. A 2016 study by the Sloan Digital Sky Survey suggested that the number is likely to be close to the lower end of that estimate, at 100–140 billion stars.The Milky Way has several components: a disk, in which the Sun and its planetary system are located; a central bulge; and a halo of stars, globular clusters, and diffuse gas. The disk is the brightest part of theMilky Way, as seen from Earth. It has a spiral structure with dusty arms. The disk is about 100,000 light-years in diameter and about 13 billion years old. It contains the young and relatively bright population I stars, as well as intermediate-age and old stars of population II.The galactic bulge is a tightly packed group of mostly old stars in the center of the Milky Way. It is estimated to contain tens of billions of stars and has a diameter of about 10,000 light-years. The Milky Way's central bulge is shaped like a box or peanut. The galactic center, which lies within this bulge, is an extremely active region, with intense radio source known as Sagittarius A*, which is likely to be a supermassive black hole.The Milky Way's halo is a spherical component of the galaxy that extends outward from the galactic disk, as far as 200,000 light-years from the galactic center. It is relatively sparse, with only about one star per cubic parsec on average. The halo contains old population II stars, as well as extremely old globular clusters.The Milky Way's spiral structure is uncertain, and there is currently no consensus on the nature of the Milky Way's spiral arms. Different studies have led to different results, and it is unclear whether the Milky Way has two, four, or more spiral arms. The Milky Way's spiral structure is thought to be a major feature of its disk, and it may berelated to the generation of interstellar matter and star formation.The Milky Way's spiral arms are regions of the disk in which the density of stars, interstellar gas, and dust is slightly higher than average. The arms are thought to be density waves that spiral around the galactic center. As material enters an arm, the increased density causes the material to accumulate, thus causing star formation. As the material leaves the arm, star formation decreases.The Milky Way's spiral arms were first identified in the 1950s, when radio astronomers mapped the distribution of gas in the Milky Way and found that it was concentrated in spiral patterns. Since then, astronomers have used a variety of techniques to study the Milky Way's spiral structure, including observations of the distribution of young stars, star-forming regions, and interstellar gas and dust.One of the key challenges in studying the Milky Way's spiral structure is that we are located within the disk of the galaxy, which makes it difficult to get a clear view of the overall structure. Astronomers have had to rely on indirect methods, such as measuring the distances and motions of stars and gas clouds, to infer the shape and structure of the galaxy.Despite these challenges, our understanding of the Milky Way's spiral structure has advanced significantly in recent years, thanks tonew observations and more sophisticated modeling techniques. Ongoing research is continuing to shed light on the nature and evolution of the Milky Way's spiral arms, and the role they play in the overall structure and dynamics of the galaxy.。

a r X i v :a s t r o -p h /9707060v 1 4 J u l 1997METAL ABUNDANCES OF ONE HUNDRED HIPPARCOS DW ARFSR.G.Gratton 1,E.Carretta 2,G.Clementini 2,C.Sneden 31Osservatorio Astronomico di Padova,Vicolo dell’Osservatorio 5,35122Padova,ITALY2Osservatorio Astronomico di Bologna,ITALY3Department of Astronomy,The University of Texas at AustinABSTRACTAbundances for Fe,O,and the α−elements (Mg,Si,Ca,and Ti)have been derived from high resolution spectra of a sample of about one hundred dwarfs with high precision parallaxes measured by HIPPAR-COS.The stars have metal abundances in the range −2.5<[Fe/H]<0.2.The observational data set con-sists of high dispersion (20,000<R <70,000),high S/N (>200)spectra collected at the Asiago and McDonald Observatories.The abundance analysis followed the same precepts used by Gratton et al.(1997a)for ∼300field stars and for giants in 24glob-ular clusters (Carretta &Gratton 1997),and includes corrections for departures from LTE in the formation of O lines.Our main results are:1.the equilibrium of ionization of Fe is well satisfied in late F –early K dwarfs2.O and α−elements are overabundant by ∼0.3dex This large homogeneous data set was used in the derivation of accurate ages for globular clusters (See paper by Gratton et al.at this same Meeting).Key words:Stars:chemical abundances -Stars:ba-sic parameters1.INTRODUCTIONHIPPARCOS has provided parallaxes with accura-cies of ∼1mas for several hundreds dwarfs.We had access to data for about 100dwarfs with metal abun-dances in the range −2.5<[Fe/H]<0.2and have used them in a thorough revision of the ages of the old-est globular clusters derived by Main Sequence (MS)fitting technique.A crucial step in the derivation of ages via this method is the assumption that the nearby subdwarfs have the same chemical composi-tion of the globular cluster main sequence stars.This assumption was verified through a careful abundance analysis of the vast majority of nearby dwarfs with HIPPARCOS parallaxes available to us.Our data set and the HIPPARCOS parallaxes were also used to test whether an appreciable Fe overion-ization occurred in the atmosphere of late F –early K dwarfs (Bikmaev et al.1990;Magain &Zhao 1996).This was done by comparing abundances provided by neutral and singly ionized lines,once the surface gravity of each program star had be derived from its mass,temperature and luminosity rather then from the equilibrium of ionization of Fe.Finally,our abundances are fully consistent with those presented by Gratton et al.(1997a)for about 300field dwarfs.A large,homogenous data base of high accuracy (errors ∼0.07dex)abundances com-puted with the Kurucz (1993)model atmospheres is now available and can be used to recalibrate photo-metric and low S/N spectroscopic abundances.2.BASIC DATA FOR SUBDWARFS Average V magnitudes and colors (Johnson B −V and V −K ,and Str¨o mgren b −y ,m 1and c 1)for the programme stars were obtained from a careful discussion of the literature data.We used also the Tycho V magnitudes and B −V colors,after cor-recting them for the very small systematic difference with ground-based data.Absolute magnitudes M V were derived combining ap-parent V magnitudes and Hipparcos parallaxes.No Lutz-Kelker corrections were applied.Lutz-Kelker corrections (Lutz &Kelker 1973)take into account that stars with parallaxes measured too high are more likely to be included in a sample if the sample selection criteria are based on the parallaxes them-selves.Since our sample was selected before the HIP-PARCOS parallaxes were known;Lutz-Kelker correc-tions should not be applied when the whole sample is considered,as we do when comparing the abundances obtained from Fe I and Fe II lines.Multiple high precision radial velocity observations exist for a large fraction of our objects (80out of 99).Twenty stars in the sample are known and four are suspected spectroscopic binaries.Two further stars display very broad lines in our spectra,possibly due to fast rotation.They were discarded.A few other stars display some IR excess,which also may be a signature of binarity.No evidence for binarity dis-turbing the present analysis exists for the remaining stars.Sixty-eight out of the99stars of our sample are in-cluded in Carney et al.(1994)catalogue.Reddening estimates are given for58of them.All but two have zero values.We have thus assumed a zero reddening for all the programme stars.3.OBSERVATIONS AND REDUCTIONSHigh dispersion spectra for about two thirds of the programme stars were acquired using the2D-coud`e spectrograph of the2.7m telescope at McDonald Ob-servatory and the REOSC echelle spectrograph at the 1.8m telescope at Cima Ekar(Asiago).McDonald spectra have a resolution R=70,000,S/N∼200, and spectral coverage from about4,000to9,000˚A; they are available for21stars(most with[Fe/H]<−0.8).Cima Ekar telescope provided spectra with resolution R=15,000,S/N∼200,and two spectral ranges(4,500<λ<7,000and5,500<λ<8,000˚A) for65stars.Equivalent widths EW s of the lines were measured by means of a gaussianfitting routine applied to the core of the lines;appropriate average corrections were included to take into account the contribution of the damping wings.Only lines with log EW/λ<−4.7 were used in thefinal analysis(corrections to the EW s for these lines are≤7m˚A,that is well be-low10per cent).The large overlap between the two samples(14stars)allowed us to tie the Asiago EW s to the McDonald ones.External checks on our EW s are possible with Ed-vardsson et al(1993:hereinafter E93)and Tomkin et al.(1992:hereinafter TLLS).Comparisons per-formed using McDonald EW s alone show that they have errors of±4m˚A.From the r.m.s.scatter,σ, between Asiago and McDonald EW s,we estimate that the former have errors of±6.7m˚A.When Asi-ago and McDonald EW s are considered together,we find average residuals(us-others)of−0.2±1.0m˚A (39lines,σ=6.1m˚A)and+0.8±1.0m˚A(36lines,σ=5.9m˚A)with E93and TLLS,respectively.4.ANALYSIS4.1.Atmospheric ParametersThe abundance derivation followed precepts very similar to the reanalysis of∼300field and∼150 globular cluster stars described in Gratton et al. (1997a)and Carretta&Gratton(1997).The same line parameters were adopted.The effective tem-peratures were derived from B−V,b−y,and V−K colours using the iterative procedure outlined in Gratton et al.(1997a).Atmospheric parameters are derived as follows:1.we assume as input values log g=4.5and themetal abundance derived from the uvby photom-etry using the calibration of Schuster&Nissen (1989)2.T effis then derived from the colours,using theempirical calibration of Gratton et al.(1997a) for population I stars(assumed to be valid for [Fe/H]=0),and the abundance dependence given by Kurucz(1993)models3.afirst iteration value of log g is then derived fromthe absolute bolometric magnitude(derived from the apparent V magnitude,parallaxes from Hip-parcos,and bolometric corrections BC from Ku-rucz1993),and masses obtained by interpolation in T effand[A/H]within the Bertelli et al.(1997) isochrones4.steps2and3are iterated until a consistent setof values is obtained for T eff,log g,and[A/H] 5.the EW s are then analyzed,providing new val-ues for v t and[A/H](assumed to be equal to [Fe/H]obtained from neutral lines)6.the procedure is iterated until a new consistentset of parameters is obtained4.2.Error analysisRandom errors in T eff(±45K)were obtained by com-paring temperatures derived from different colours. Systematic errors may be larger;the T eff-scale used in this paper is discussed in detail in Gratton et al. (1997a).We assume that systematic errors in the adopted T eff’s are≤100K.Random errors in the gravities(±0.09dex)are esti-mated from the errors in the masses(1.2per cent), M V’s(0.18mag),and in the T eff’s(0.8per cent), neglecting the small contribution due to BC’s.Sys-tematic errors(±0.04dex)are mainly due to errors in the T effscale and in the solar M V value.Random errors in the microturbulent velocities can be estimated from the residuals around thefitting re-lation in T effand log g.We obtain values of0.47and 0.17km s−1for the Asiago and McDonald spectra, respectively.Random errors in the EW s and the line parameters significantly affect the abundances when few lines are measured for a given specie.Errors should scale as σ/√Figure1.Run of the difference between the abundances derived from neutral and singly ionized Fe lines as a func-tion of temperature(panel a)and overall metal abundance (panel b).Open squares are abundances obtained from the Asiago spectra;filled squares are abundances obtained from the McDonald spectraMcDonald spectra,respectively.Systematic errors (∼0.08dex)are mainly due to the T effscale.parison with other abundancesOn average,differences(Asiago−McDonald)in the Fe abundances are−0.01±0.02dex(12stars,σ= 0.07dex).Analogous differences for the[O/Fe] and[α/Fe]ratios are+0.02±0.08dex(5stars,σ=0.17dex),and+0.01±0.03dex(12stars,σ=0.10dex).E93measured abundances for∼200dwarfs;six stars are in common with our sample.Abundance residu-als(our analysis−E93)are+0.08±0.03,−0.02±0.03, and+0.02±0.02dex for[Fe/H],[O/Fe],and[α/Fe], respectively.Residual differences are mainly due to our use of a higher temperature scale(our T eff’s are larger by63±12K).We have six stars in common with TLLS,which used a restricted wave-length range.Average differences(ours−TLLS)are: +0.34±0.04and−0.31±0.07dex for[Fe/H]and [O/Fe],respectively.They are due to different as-sumption in the analysis:(i)our temperature scale is higher;(ii)TLLS used a different solar model; (iii)our non-LTE corrections to the O abundances are slightly larger.Finally,Gratton et al.(1997a) made a homogenous reanalysis of the original EW s for∼300metal-poorfield stars.On average,the present Fe abundances are larger by0.02±0.02dex (11stars,σ=0.06dex).Since the same analysis procedure is adopted,these differences are entirely due to random errors in the EW s and in the adopted colours.In the following,we assume that Gratton et al.abundances are on the same scale of the present analysis.4.4.Fe abundancesSince gravities are derived from masses and luminosi-ties rather than from the equilibrium of ionization for Fe,we may test if predictions based on LTE are sat-isfied for the program stars.In Figure1we plot the difference between abun-dances of Fe obtained from neutral and singly ion-ized lines against effective temperature and metal abundance.Different symbols refer to results ob-tained from McDonald and Asiago spectra,respec-tively.McDonald spectra have a higher weight be-cause the higher resolution allowed us to measure a larger number of Fe II lines(10∼20),and errors in the EW s are smaller;very few Fe II lines could be measured in the crowded spectra of cool and/or metal-rich stars observed from Asiago.Average dif-ferences between abundances given by Fe I and II lines are0.025±0.020(21stars,σ=0.093dex)for the Mc Donald spectra,and−0.063±0.019(52stars,σ=0.140dex)for the Asiago spectra.The scatter obtained for McDonald spectra agrees quite well with the expected random error of0.085dex.The average value is consistent with LTE if the adopted T effscale is too high by∼20K,well within the quoted error bar of±100K.The lower mean difference obtained for the Asiago spectra is due to a few cool metal-rich stars which have very crowded spectra.Very few Fe II lines could be measured in these spectra and the line-to-line comparison with the superior McDonald data suggests that even these lines may be affected by blends.We conclude that the equilibrium of ionization for Fe is well satisfied in the late F–K dwarfs of any metallicity in our sample.This result depends on the adopted temperature scale.Our empirical result agrees very well with the ex-tensive statistical equilibrium calculations for Fe by Gratton et al.(1997b).In that paper,the uncertain collisional cross sections were normalized in order to reproduce the observations of the RR Lyraes,where overionization is expected to be much larger than in late F–K dwarfs.The lower limit to collisional cross sections given by the absence of detectable overion-ization in RR Lyrae spectra(Clementini et al.1995) implies that LTE is a very good approximation for the formation of Fe lines in dwarfs.4.5.O andα−element abundancesO abundances were derived from the permitted IR triplet,and include non-LTE corrections computed for each line in each star following the precepts of Gratton et al.(1997b).Wefind that O and the other α−elements are overabundant in stars with[Fe/H]<−0.5(see Figure2):[O/Fe]=0.38±0.13[α/Fe]=0.26±0.08,(error bars are the r.m.s.scatter of individual val-ues around the mean).The moderate O excess de-rived from the IR permitted lines is a consequence of the rather high temperature scale adopted.When this adoption is made,abundances from permitted OI lines agree with those determined from the forbidden [OI]and the OH lines.The present abundances agree very well with those derived in Gratton et al.(1997c).Note also that the overabundance of O andα−elements found for thefield subdwarfs is similar to the excesses foundFigure2.Runs of the overabundances of O(panel a)and α−elements(panel b)as a function of[Fe/H]for the pro-gramme subdwarfs.Filled squares are abundances from McDonald spectra;open squares are abundances from Asi-ago spectrafor globular cluster giants(apart from those stars af-fected by the O-Na anticorrelation,see Kraft1994).5.CALIBRATION OF PHOTOMETRICABUNDANCESOnce combined with the abundances obtained by Gratton et al.(1997a),the sample of late F to early K-typefield stars with homogenous and accu-rate high dispersion abundances adds up to nearly 400stars.Schuster&Nissen(1989)have shown that rather accurate metal abundances for late F to early K-type can be obtained using Str¨o mgren uvby pho-tometry(available for a considerable fraction of the HIPPARCOS stars).Furthermore,the extensive bi-nary search by Carney et al.(1994)has provided a large number of metal abundances derived from an empirical calibration of the cross correlation dips for metal-poor dwarfs.We have recalibrated these abundance scales.Schus-ter&Nissen(1989)abundances onlydiffers for a zero-point offset(see panel a of Figure3);the mean difference is:[Fe/H]us =[Fe/H]SN+(0.102±0.012),(1)based on152stars(the r.m.s.scatter for a single star is0.151dex).In the case of Carney et al.(1994,panel b of Fig-ure3),a small linear term is also required.The best parison between the abundances obtained from high dispersion spectra(present analysis or Gratton et al.1997),and those provided by the original calibration of Schuster&Nissen(1989,panel a)and Carney et al. (1994,panel b)fit line(66stars)is:[Fe/H]us=(0.94±0.03)[Fe/H]C94+(0.18±0.17),(2) The offsets between the high dispersion abundances and those provided by Schuster&Nissen(1989)and Carney et al.(1994)are mainly due to different as-sumptions about the solar abundances in the high dispersion analyses originally used in the calibrations of Schuster&Nissen(1989)and Carney et al.(1994).REFERENCESBertelli,P.,Girardi,L.,Bressan,A.,Chiosi,C.,&Nasi,E.1997,in preparationBikmaev,I.F.,Bobritskij,S.S.,El’kin,V.G.,Lyashko,D.A.,Mashonkina,L.I.,&Sakhibullin,N.A.1990,inIAU Symp.145,Evolution of Stars:the Photospheric Abundance Connection,G.Michaud ed.Carney,B.W.,Latham,D.W.,Laird,J.B.,&Aguilar, L.A.1994,AJ,107,2240Carretta,E.,&Gratton,R.G.1997,A&AS,121,95 Clementini,G.,Carretta,E.,Gratton,R.G.,Merighi, R.,Mould,J.R.,&McCarthy,J.K.1995,AJ,110, 2319Edvardsson,B.,Andersen,J.,Gustafsson,B.,Lambert,D.L.,Nissen,P.E.,&Tomkin,J.1993,A&A,275,101Gratton,R.G.,Carretta,E.,&Castelli,F.1997a,A&A, in pressGratton,R.G.,Carretta,E.,Gustafsson,B.,&Eriksson, K.1997b,submitted to A&AGratton,R.G.,Carretta,E.,Matteucci,F.,&Sneden,C.1997d in preparationKing,J.R.1993,AJ,106,1206Kraft,R.P.1994,PASP,106,553Kurucz,R.L.1993,CD-ROM13and CD-ROM18 Lutz,T.E.,Kelker,D.H.1973,PASP,85,573 Magain,P.,Zhao,G.1996,A&A,305,245Schuster,W.J.,&Nissen,P.E.1989,A&A,221,65 Tomkin,J.,Lemke,M.,Lambert,D.L.,&Sneden,C.1992,AJ,104,1568。