Chemical Analysis of the Fornax Dwarf Galaxy

University of Groningen

Chemical analysis of the Fornax dwarf galaxyLetarte, Bruno

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Rĳksuniversiteit Groningen

Chemical Analysis of theFornax Dwarf Galaxy

Proefschrift

ter verkrĳging van het doctoraat in deWiskunde en Natuurwetenschappenaan de Rĳksuniversiteit Groningen

op gezag van deRector Magnificus, dr. F. Zwarts,in het openbaar te verdedigen op

vrĳdag 30 maart 2007om 14.45 uur

door

Bruno Letarte

geboren op 12 juni 1976te Québec, Canada

Promotor: Prof. dr. E. TolstoyCopromotor: Dr. V. Hill

Beoordelingscommissie: Prof. dr. M. SpiteProf. dr. P. C. van der KruitProf. dr. J. W. Pel

ISBN 90-367-2927-0ISBN 90-367-2928-9 (electronic version)

In the beginning the Universe wascreated. This has made a lot of peo-ple very angry and has been widelyregarded as a bad move.

–Douglas Adams

Cover page – Fornax Dwarves, by Jesse Giroux

Contact information:

Bruno [email protected]@astro.rug.nl

This thesis has been funded by:

With support from:

LKBFLeids Kerkhoven-Bosscha Fonds

Contents

1 Introduction 91.1 The Cosmological Importance of Dwarf Galaxies . . . . . . . . . . . . . . 91.2 The Formation of the Elements . . . . . . . . . . . . . . . . . . . . . . . . 111.3 Abundances in Galaxies . . . . . . . . . . . . . . . . . . . . . . . . . . . . 12

1.3.1 The Milky Way . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 131.3.2 The Magellanic Clouds & Dwarf Galaxies . . . . . . . . . . . . . . 14

1.4 The DART project . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 151.4.1 Photometry . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 151.4.2 Spectroscopy . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 151.4.3 This Thesis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17

2 Fornax and the Local Group 192.1 Dwarf galaxies in the Local Group . . . . . . . . . . . . . . . . . . . . . . 192.2 Fornax dSph . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 212.3 Globular Clusters in Fornax . . . . . . . . . . . . . . . . . . . . . . . . . . 24

3 Using stellar atmospheric models ... chemical abundances 273.1 Describing the stellar atmosphere . . . . . . . . . . . . . . . . . . . . . . . 27

3.1.1 The flux . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 283.1.2 The absorption coefficient . . . . . . . . . . . . . . . . . . . . . . . 303.1.3 Stellar atmospheric models . . . . . . . . . . . . . . . . . . . . . . 35

3.2 Determining Stellar Atmospheric parameters . . . . . . . . . . . . . . . . 353.2.1 Effective Temperature (Teff) . . . . . . . . . . . . . . . . . . . . . . 353.2.2 Surface Gravity (log g) . . . . . . . . . . . . . . . . . . . . . . . . . 363.2.3 Metallicity . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 393.2.4 Microturbulence velocity . . . . . . . . . . . . . . . . . . . . . . . . 39

3.3 The abundance determination . . . . . . . . . . . . . . . . . . . . . . . . . 393.3.1 Measuring the equivalent widths . . . . . . . . . . . . . . . . . . . 403.3.2 The Stellar Models used . . . . . . . . . . . . . . . . . . . . . . . . 403.3.3 Computing the abundances . . . . . . . . . . . . . . . . . . . . . . 42

3.4 The line list . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 433.4.1 Building a line list . . . . . . . . . . . . . . . . . . . . . . . . . . . 43

vi CONTENTS

3.4.2 The line by line selection . . . . . . . . . . . . . . . . . . . . . . . 45

4 Abundances with the FLAMES multi-fibre instrument 474.1 UVES vs FLAMES . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 48

4.1.1 UVES . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 494.1.2 FLAMES . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 49

4.2 The FLAMES Spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 504.2.1 Extracting, calibrating . . . . . . . . . . . . . . . . . . . . . . . . . 504.2.2 Combining . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 504.2.3 Determining the radial velocities (Vrad) . . . . . . . . . . . . . . . 524.2.4 Measuring the Equivalent Widths . . . . . . . . . . . . . . . . . . 524.2.5 Cleaning up the spectra . . . . . . . . . . . . . . . . . . . . . . . . 56

4.3 Selecting our stellar parameters . . . . . . . . . . . . . . . . . . . . . . . . 574.3.1 Photometric gravity . . . . . . . . . . . . . . . . . . . . . . . . . . 574.3.2 Photometric Teff . . . . . . . . . . . . . . . . . . . . . . . . . . . . 574.3.3 Iterating on the parameters . . . . . . . . . . . . . . . . . . . . . . 624.3.4 Precision and error estimates . . . . . . . . . . . . . . . . . . . . . 65

4.4 Systematics and corrections . . . . . . . . . . . . . . . . . . . . . . . . . . 684.4.1 Systematics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 684.4.2 Hyperfine splitting correction . . . . . . . . . . . . . . . . . . . . . 71

Appendix 4.A Large tables . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 72

5 HR spectroscopy in Fornax Globular Clusters 775.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 785.2 Observations . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 795.3 Data Reduction and Analysis . . . . . . . . . . . . . . . . . . . . . . . . . 815.4 Interpretation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 85

5.4.1 The Iron abundance . . . . . . . . . . . . . . . . . . . . . . . . . . 855.4.2 The Alpha elements . . . . . . . . . . . . . . . . . . . . . . . . . . 865.4.3 Deep mixing pattern . . . . . . . . . . . . . . . . . . . . . . . . . . 895.4.4 Iron-peak elements . . . . . . . . . . . . . . . . . . . . . . . . . . . 905.4.5 Heavy elements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 92

5.5 Conclusions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 95Appendix 5.A Large tables . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 97

6 HR spectroscopic study of Fornax Field Stars 1056.1 Sample selection . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1066.2 Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 107

6.2.1 Iron abundance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1076.2.2 Alpha Elements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1086.2.3 Iron peak elements . . . . . . . . . . . . . . . . . . . . . . . . . . . 1136.2.4 Deep-mixing pattern . . . . . . . . . . . . . . . . . . . . . . . . . . 1146.2.5 The Na-Ni relationship . . . . . . . . . . . . . . . . . . . . . . . . 1156.2.6 Heavy elements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 116

6.3 Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1206.3.1 Comparison of Fornax and Sculptor . . . . . . . . . . . . . . . . . 1216.3.2 Age and [Fe/H] . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 122

6.4 Conclusions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 123

CONTENTS vii

Appendix 6.A Large tables . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 124

7 Conclusions 1417.1 New Data Reduction and Analysis Techniques . . . . . . . . . . . . . . . . 1417.2 The Fornax Globular Clusters . . . . . . . . . . . . . . . . . . . . . . . . . 1427.3 Fornax Field stars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 142

Bibliography 145

Nederlandse samenvatting 151

Résumé français 155

Acknowledgements 159

Chapter 1Introduction

Dwarf galaxies are in principle the most simple and straightforward type of galaxyand their study can be used to test numerous theories of the formation and evolution

of stars and galaxies in a range of environments. This thesis concentrates on the detailedstudy of the chemical elements in individual stars in the nearby dwarf spheroidal galaxy,Fornax. A dwarf spheroidal galaxies are small roughly spherical galaxies that are typicallyfound in the vicinity of larger galaxies, such as the Milky Way. They typically do not haveany ongoing star formation, nor to they appear to have any gas associated to them. Theabundance ratios of different elements in individual stars with a range of ages provide adetailed insight into the various chemical enrichment processes (e.g., supernovae, stellarwinds) which in turn improves our understanding of the global processes of formationand evolution of a galaxy as a whole.

1.1 The Cosmological Importance of Dwarf GalaxiesThe most straightforward model of galaxy formation is that all galaxies form in the earlyUniverse in a rapid collapse scenario (so called monolithic collapse, Eggen, Lynden-Bell,& Sandage 1962). These galaxies then evolve solely by changing their gas mass into astellar mass with time. This model assumes that the majority of the mass of all galaxieswas in place at their formation. However this basic picture was updated (e.g., Searle &Zinn 1978) to a model which assumes that galaxies are not formed in a single collapse,but that they are built up in time from smaller fragments. This theory came in parallelwith the very successful “cold dark matter” (CDM) vision of structure formation in theUniverse which assumes that the dark matter content of a galaxy is built up through thecontinuous accretion of small clumps, to build up the galaxies and clusters of galaxies wesee today (e.g., White & Rees 1978; Navarro, Frenk, & White 1995).

If we take the CDM model of structure formation and assume that the ratio of bary-onic to dark matter is roughly constant and known then this naturally results in theconcept of numerous “building blocks”, or small galaxies, which are continuously beingaccreted onto larger galaxies over the history of the Universe. These small galaxies, with

10 chapter 1: Introduction

a similar mass to the dwarf galaxies we see today, might act as stellar nurseries, creatingthe stars we see in the Milky Way (MW) today. Stars within the Galactic halo are someof the oldest objects ever observed and they should be representative of the earliest starformation in the Local Group (LG). These stars either formed in the proto-Milky Way orthey may have formed in smaller satellite galaxies that were accreted to the Milky Wayat a later time. CDM based models thus suggest that a considerable fraction of the starsin the Milky Way today should have formed in smaller building blocks. For example, theSagittarius dwarf galaxy behaves exactly like a CDM building block, showing signs ofbeing tidally disrupted and merging in its entirety into the Milky Way (Ibata et al. 1994).

As required by the CDM view of the Universe small galaxies do appear to be darkmatter dominated (e.g., Mateo 1998). Observations of dwarf spheroidal galaxies in theLocal Group, such as Fornax dSph, suggest that dwarf galaxies must be considerablymore massive than the visible mass would suggest (e.g., ∼ 109 − 1010M�, as comparedto visible masses of ∼ 107 − 108M�), (Mateo et al. 1991; Walker et al. 2006; Battagliaet al. 2006). However there are inconsistencies in the predicted properties of the DMprofiles of the observed dwarfs and the predictions of CDM (e.g., Wilkinson et al. 2006).It also appears that the properties of the stellar populations, the dark to baryonic matterratio, and the kinematic properties of dwarf galaxies we see today are inconsistent withthe requirements of building blocks of the Milky Way, i.e., adding together all the smallgalaxies we see today, or at any time in the past, will not result in a galaxy like the MilkyWay (e.g., Shetrone et al. 2003; Tolstoy et al. 2003; Venn et al. 2006; Helmi et al. 2006).

CDM also appears to over-predict the number of small satellite galaxies around largergalaxies such as our own, an inconsistency that is known as the “missing dwarf problem”(e.g., Moore et al. 1999). However, recent discoveries of several faint satellites aroundthe Milky Way in the last couple of years are changing our view about the LG (e.g.,Belokurov et al. 2006a,b; Willman et al. 2005a,b; Zucker et al. 2006a,b). These studiessuggest that the dwarf spheroidal galaxies we have studied to date are only the tip of theiceberg; they are the most massive satellites of a larger population of fainter, lower masssatellites (e.g., Stoehr et al. 2002), which could bring our Milky Way environment backinto consistency with the CDM predictions for the amount of sub-structure. However,our knowledge about these new faint galaxies, especially their dark matter content, isstill quite limited as they have been discovered relatively recently.

Thus dwarf galaxies are useful probes of our understanding of galaxy formation andevolution on the smallest scales and potentially also as building blocks of the largestgalaxies. By studying the nearest examples we can obtain the kind of detailed com-parisons between theory and observation that are required to test current theories andprovide a solid observational basis for future models. More specifically, studying theabundance patterns of stars of a range of age allows us to understand in detail the evo-lutionary processes that shape galaxies everywhere. Thus looking at individual stars indwarf galaxies in the Local Group is an important component in understanding the bigpicture of galaxy formation and evolution throughout the Universe.

1.2: The Formation of the Elements 11

1.2 The Formation of the Elements

It is believed that the Universe started as an explosion, known as the Big Bang, wherehydrogen, deuterium, helium and lithium were created. These are thus considered to beprimordial elements and all other elements are formed subsequently by nucleosynthesisin stars. Stellar nucleosynthesis is thus responsible for almost all of what we see aroundus on the Earth today. It was first explained in the 1950s in work done by Fowler andHoyle, culminating in the B2FH (Burbidge, Burbidge, Fowler, & Hoyle 1957) paper.

The first most fundamental process of converting hydrogen into heavier elements ishydrogen burning, which is the conversion of hydrogen nuclei into helium, via the proton-proton chain in low mass stars with low core temperatures, and via proton captures bycarbon, nitrogen, and oxygen atoms (in the CNO cycles) in more massive stars withhigher temperatures. The CNO cycle traces the origin of most of the observed nitrogentoday, while most of the helium produced is consumed in the next stage: helium burning.As helium builds up in the core of the star the core contracts until the temperature anddensity increase enough to allow for another reaction in which helium is the fuel. Thisthermonuclear phase is the triple-α process in which three 4He nuclei fuse to form acarbon nucleus. The next stage is shell burning: carbon burning, oxygen burning, siliconburning. This can produce elements as heavy as 56Fe which is the most massive elementthat can be formed by fusion in the core of a star.

The most significant group of heavier elements are the so called alpha elements, withnuclei that are multiples of He, e.g., O, Mg, Ca, Si and Ti. They are predominantlysynthesised by alpha capture during the various burning phases in massive stars, andexpelled into the ISM by SN II explosions. Another significant and important group ofelements is the iron-peak, including Fe itself. It is predominantly produced and expelledinto the ISM by SN Ia, supernovae thought to be due to the explosion of a white dwarfin an evolved binary with a less massive progenitor star. Those typically occur ∼1 Gyrafter the first episode of star formation, contrary to SN II which have short-lived massivestar progenitors (as short as ∼10 Myrs). As a consequence, elemental ratios of the type[α/Fe] inform us of the relative contribution from the two types of supernovae at a giventime, indicative of star formation timescale. Figure 1.1 sketches how the [α/Fe] ratio canbe viewed as a kind of chronometer (starting to decrease after 1Gyr) while the [Fe/H]metallicity index provides the efficiency with which star formation has occurred. Whenthe star formation rate (SFR) is high, then the gas will reach higher [Fe/H] before thefirst SN Ia occur and α-elements start to decrease (the “knee”). The formation efficiencyand time scale of a stellar system can be estimated by the position of this “knee”. And,because more massive stars are more efficient in producing α-elements, the level of [α/Fe]at low metallicity (before the “knee”) is an indication of the mass of the stars that con-tributed to enrich the ISM and therefore provides a indirect measure of the IMF.

Heavier elements beyond the iron peak are created by neutron capture, where the twomost important processes (in the astrophysical context) are the s- and r- processes. Thes-process (or slow-process) occurs when the neutron flux is not very high, so that theintervals between neutron captures are long compared to the beta decay characteristictimescale of an unstable nucleus. These conditions are found in the envelopes of ther-


Figure 1.1: Simpleview of how α-elementscan be used to tracethe IMF and SFH ofa galaxy (taken fromMcWilliam 1997).

mally pulsating AGB stars, and are most efficient in 3-5M� stars. Because of the slowevolution of intermediate-mass stars, s- process will only enter the chemical enrichmentof a galaxy several 100 Myrs after the first episode of star formation. In addition, itrequires pre-existing iron-peak elements seeds in the AGB envelope, and is therefore in-efficient at very low metallicity. The s- process is unlikely to be significant in the earlieststages of star formation in a galaxy.

The r-process (or rapid-process) occurs when there is sufficient neutron flux whichallows rapid captures of neutrons. This is believed to occur predominantly in environ-ments like those produced by SNe II. With such rapid successive captures, neutrons canaccumulate on an unstable nuclei before it has time to either beta or alpha decay. Thestars responsible for these explosions are massive, therefore have a short lifetime and arebelieved to be the first objects that will contribute heavy elements to the ISM. Observingthe relative abundances of s- and r- process nuclei can therefore constrain the impact ofAGB stars on chemical evolution and probe star formation timescales.

1.3 Abundances in Galaxies

Because elemental abundances are preserved∗ at the stellar surface during the wholestellar lifetime, and can be (relatively) easily measured from absorption lines in high-resolution stellar spectra, they have become a very important tool to understand thegenesis of a stellar population. Abundances of various elements can be measured in starsof different ages and, thanks to their different nucleosynthetic origin, allow us to inferwhat enrichment processes have been dominant at different epochs of galaxy formation.Not surprisingly our earliest studies have concentrated on the Milky Way, and it is onlyrelatively recently that similarly detailed studies have been made of other galaxies, suchas the Magellanic Clouds and most recently the nearby dwarf spheroidal galaxies.

∗ Except for a few light elements which may be affected by internal mixing: Li, C, N.

1.3: Abundances in Galaxies 13

1.3.1 The Milky WayThe Milky Way contains several stellar components which are distinguished by differ-ent spatial distribution, kinematics and stellar populations, namely the halo, the thickdisk, the thin disk, and the bulge. Each component has clearly had a different forma-tion history and their stars show marked differences in their age distribution, metallicitydistribution and most importantly here, abundance ratios. Ever since the discovery byChamberlain & Aller (1951) that two stars with high radial velocities (halo stars) hadtheir iron and calcium abundances an order of magnitude lower than that of the Sun,it gradually became clear that the various stellar populations that comprise the MilkyWay have both kinematics and chemical signatures associated to each of them. and thatcombining the two properties was necessary to better understand galaxy evolution (Wyse& Gilmore 1995).

A review by McWilliam (1997), covering the Galactic disk, halo and bulge suggestthat the environment plays an important role in chemical evolution and that supernovaecome in many flavors, with a range of element yields. Below are a some recent examplesof detailed abundance studies of the Milky Way:

The detailed abundance studies of extremely metal poor stars in the halo of ourGalaxy have given us a clearer picture of its earliest enrichment history. The high [Zn/Fe]observed and absence of very strong depletion of odd-numbered elements have ruled outpair instability SN (from 130-300M� progenitors) as a dominant source of enrichment(Cayrel et al. 2004). The dispersion in heavy neutron-capture element abundances of themost metal poor stars suggests incomplete mixing of the ejecta from individual super-novae into the galactic interstellar medium (McWilliam 1997).

Studies of large samples (∼200) nearby disk stars (F and G dwarf) provide observa-tional constraints by linking chemical abundance of up to 30 chemical elements to precisekinematics and photometric ages (e.g., Edvardsson et al. 1993; Chen et al. 2000; Reddyet al. 2003). This has allowed to understand that the thin disk formed stars at a steadyrate over the last 4-8 Gyrs, allowing a full evolution of the abundance ratios from almostpure SN II ejecta to a full mix of SN II, stellar winds and SN Ia. Although the meanmetallicity increases with time, the age-metallicity relation is neither well defined nortight in the galactic disk, ruling out the “instantaneous mixing” assumption of simplemodels of galaxy chemical evolution.

Recent precision work has shown that the [α/Fe] ratio for thick-disc stars shows aclear enhancement compared to thin-disc members of the same metallicity, which is asign that star formation was more efficient and restricted to a shorter period of time inthe thick disk (e.g. Bensby et al. 2003, 2005; Reddy et al. 2006). Several hypotheses havebeen proposed for the origin of the thick disk: the debris of a merger, a merger thatheated a preexisting thin disk into a thick disk, etc. The first indications of a populationthat could be ascribed to debris from the satellite whose merger caused the thick diskwas presented in Gilmore et al. (2002): thick disk stars should then bear the chemicalsignature of the star formation history of the merging (dwarf ?) galaxy. Galactic starsseen along lines of sight to some dSph galaxies seem to have the expected properties of“satellite debris” in the thick disk-halo interface, which is interpreted as remnants of the


merger that heated a preexisting thin disk to form the thick disk (Wyse et al. 2006).Thick disk stars would then have the chemical signature of the former thin disk.

In studies of Galactic bulge stars, two α-element ratios, [O/Fe] and [Mg/Fe] have beenfound to be higher than in thick disk stars, which are known to be more oxygen rich thanthin disk stars (e.g., Zoccali et al. 2006; Fulbright et al. 2006; Lecureur et al. 2006). Thissupports a scenario in which the bulge formed before and more rapidly than both thethin and thick disks, and therefore the MW bulge can be regarded as a prototypical oldspheroid, with a formation history similar to that of early-type (elliptical) galaxies.

1.3.2 The Magellanic Clouds & Dwarf GalaxiesOther galaxies are in principle simpler to interpret than the Milky Way as we have anexternal view of the entire system and distance differences are unimportant. Stars in theMagellanic Clouds (at ∼50 kpc distance) were the first extragalactic stars targeted fordetailed abundance studies and the results of these studies gave us the first insights intoa more metal poor star forming environment than is available in the disk of our galaxy.At the end of the 80’s and in the 90’s, 4m-class telescopes were used to study detailedabundances of supergiant stars in both the Large and Small Magellanic Clouds, reflectingthe current interstellar medium within these galaxies (Russell & Bessell 1989; Hill et al.1995; Hill 1997; Venn 1999). Probing the chemical composition of stars as a function ofage and therefore chemical evolution per se had to wait until 8-10m class telescopes gaveaccess to high-resolution spectra of RGB stars in the Large Magellanic Cloud, initiallyin small numbers, (Hill et al. 2000; Smith et al. 2002), followed by the first abundancestudy of a large sample (Pompeia et al. 2006).

Similarly, the Sagittarius dSph has also been targeted in high resolution studies ofsome tens of RBG stars (Bonifacio et al. 2000; Monaco et al. 2005). These studies re-vealed distinctive evolutionary paths for the Large Magellanic Cloud and the SagittariusdSph, showing a different chemical enrichment process from the Milky Way and otherdwarf galaxies (Bonifacio et al. 2000).

However the Magellanic Clouds and Sagittarius are clearly in the process of inter-acting strongly with our galaxy and so the lessons they have to teach about galaxyformation and evolution are not so straightforward to interpret. Dwarf galaxies, on theother hand, especially the nearby dSph are arguably simpler and more clearly preservedenvironments. These are however twice as distant as the Magellanic Clouds, and thusdetailed abundances require 8-10m class telescopes.

Using Keck to look at individual stars in the Draco, Sextans and Ursa Minor dSph(Shetrone et al. 1998, 2001), and soon after the VLT for four southern dSph (e.g., Shetroneet al. 2003; Tolstoy et al. 2003), studies of LG dSph were initially based on very smallsamples of stars and yet they provided fundamental insights into galaxy formation andevolution. From these studies it became evident that, whereas the metallicity of dSphstars seemed to lie between the bulk of Galactic disk and halo stars, α-elements weretypically under abundant when compared to MW stars of similar metallicity (hence lowerthan in the halo), while r- and s- process elements in dSph stars were typically halo-like.

1.4: The DART project 15

This suggests that the satellite galaxies we see today cannot be significant recentcontributors to the stellar population of our Galaxy, with the possible exception of theouter halo. However, the lack of statistically significant samples of objects (2−5 starsper dSph) undermined the strength of this conclusion.

More importantly still, although dSph are simpler systems when compared to theMW, with most of them having typically much lower star formation rates, each of themhas a unique and different star formation history. Abundance ratios were yet to be studiedin large enough samples in several different dSph to understand the internal evolution ofthese systems.

1.4 The DART projectDART is an acronym for Dwarf Abundance and Radial-velocity Team (Tolstoy et al.2004, 2006). It involves more than 16 persons, from 10 institutes in 10 different countries.The main goal of the project is to obtain detailed chemical abundances (requiring highresolution) and radial velocities (low resolution) for a large sample of stars in four nearbydSph galaxies, Sculptor, Fornax, Sextans and Carina (for which we obtained high resolu-tion spectroscopy only). The project is primarily based on two observing proposals, theESO Large Programme 171.B-0588 (PI: Tolstoy) entitled: “Dwarf galaxies: remnants ofgalaxy formation and corner stones for understanding galaxy evolution” and the MeudonGTO Programme 71.B-0641 (PI: Hill) entitled: “Star formation history of the Sculptordwarf spheroidal galaxy” which began obtaining data in August 2003.

1.4.1 PhotometryWide-field accurate photometry was needed both to select targets for our spectroscopicsurvey, and to allow a colour-magnitude diagram analysis of the global properties (meanages and metallicities) of the stellar populations in the galaxy and the underlying starformation history. Precise astrometry (better than 0.3′′) of the targets selected for spec-troscopic follow-up is also required to insure a proper placement of FLAMES fibres (1.2′′fibre entrance on the sky).

The instrument we used for our photometric survey is the wide field imager WFI,(Baade et al. 1999) on the 2.2-m MPG/ESO telescope on La Silla. The large field ofview (33′× 34′) of this instrument allowed us to efficiently map out dwarf galaxies tobeyond their tidal radius. Our photometric survey was conducted in the visible band Vand I. Figure 1.2 shows the spatial distribution of our imaging for Fornax. We have alsoplotted the low-resolution spectroscopic survey, with bigger black points representing theFLAMES LR targets (Battaglia et al. 2006).

1.4.2 SpectroscopyWe used VLT/FLAMES, described in Pasquini et al. (2002) as well as in chapter 4 of thisthesis, to carry out our spectroscopic survey. For each of the four galaxies, we obtainedone FLAMES pointing in high resolution mode, each consisting of ∼100 target stars forwhich we will obtain chemical abundances (and radial velocities). To obtain sufficient


Figure 1.2:Spatial distribu-tion of the FornaxDART imagingsample. The coor-dinates ξ and η arede-projected rectan-gular coordinates,using the centre ofFornax derived in(Battaglia et al.2006). The ellipsesare drawn at 1, 2, 3and 4 core radius,Rcore, the last onecorresponding to thetidal radius, Rtidal.

wavelength coverage for an accurate analysis of the abundances and to include a varietyof chemical elements, we used several different setups that cover different wavelengthranges. Three setups were obtained in order to perform an abundance analysis, totallingalmost 30h of observation per galaxy (see chapter 4).

On the other hand, a low resolution pointing can be obtained in about an hour, al-lowing for a greater number of stars for which we get a basic metallicity tracer (Ca IItriplet, or CaT) and a radial velocity. The Fornax study in LR consisted of 11 pointingsand 1063 targets, as illustrated in Figure 1.2 (see Battaglia et al. 2006).

The DART studies to date, as well those of other groups, (e.g., Koch et al. 2006, 2007)have shown that neither the kinematics nor the metallicities nor the spatial distributionsof dSph are easy to explain in a straightforward manner even for these smallest galax-ies. Dwarf galaxies show complex and highly specific evolutionary and metal-enrichmentprocesses. Details of these results coming from low resolution CaT spectroscopy are pre-sented in Tolstoy et al. (2004) (for Sculptor) and in Battaglia et al. (2006) for Fornax.Specifically, in Fornax, we have shown that the galaxy contains at least two morpholog-ically (concentration), chemically (metallicity) and kinematically (velocity dispersion)distinct intermediate to old components. The centre of Fornax is dominated by the moremetal-rich and kinematically cooler (and younger) component. This is the populationfrom which our high-resolution sample was drawn.

1.4: The DART project 17

Figure 1.3: Digital Sky Survey (DSS) image of the central region of Fornax (85′ x 62′or 3.4 x 2.5 kpc) with the central 25′ field identified by a big circle and the five globularclusters with smaller circles.

1.4.3 This ThesisThe main emphasis of this thesis is to determine detailed chemical abundances of in-dividual stars in the nearby Fornax dwarf spheroidal galaxy, based on high resolutionobservations with VLT/FLAMES. We have targeted stars in the central 25′ diameterregion of Fornax, as well as in three of its globular clusters. An image of Fornax is shownin Figure 1.3, where the central FLAMES field we observed in HR is identified, as wellas the location of the five globular clusters of Fornax. The goal was to make a consis-tent study of the chemical properties of a representative sample of the stellar populationof Fornax, and to make a comparison between the properties of stars in its old globu-lar clusters (GCs) and predominantly intermediate age field stars. Detailed abundanceanalysis from HR spectroscopy is necessary for the full understanding of a complicatedstar formation history, where classic colour-magnitude diagram (CMD) analysis is notsufficient to provide a definitive answer.

Although earlier studies have provided hints of the evolutionary processes in dwarfgalaxies the unparallelled multi-tasking capability of VLT/FLAMES allows us to mapout the large scale processes which are important on the scale of a dSph and also todistinguish “the weather from the climate” in these galaxies – with regard to the chemicalevolution with time.

Chapter 2Fornax and the Local Group

The cold dark matter (CDM) paradigm states that small galaxies are the buildingblocks of larger galaxies. This formation scenario is quite successful at modelling

large scale structures but it has a problem for objects the size of current day dwarf galax-ies: many more dwarfs are predicted than are actually observed in the Local Group. Thesurviving dwarf galaxies give us the opportunity to learn more about them and theirrelation to larger galaxies such as the Milky Way and M 31. By studying photometricproperties, kinematics and the detailed chemistry of individual stars in different systems,both large and small, we can hope to better understand galaxy formation and evolution.

In my thesis I carry out a detailed high resolution spectroscopic study of individualstars in a nearby dwarf galaxy: the Fornax dwarf spheroidal galaxy. In this chapter Iprovide an introduction to what is currently known about Fornax and the environmentin which it is evolving.

2.1 Dwarf galaxies in the Local GroupThe Local Group contains ∼40 dwarf galaxies, mostly clustered around two big spiralgalaxies, the Milky Way (MW) and M 31 (van den Bergh 2000), see Figure 2.1 for aschematic overview. The majority of dwarf galaxies generally fall into two categories,Dwarf Spheroidals (dSphs) and Dwarf Irregulars (dIrrs). The dSphs are generally foundclose to a host galaxy and they typically don’t have current star formation or H i gasassociated with them and the dIrrs are typically more distant and generally have at leastsome current star formation and gas (Mateo 1998).

The Local Group is a useful laboratory to study galaxies in detail because, as opposedto high redshift surveys, we can resolve individual stars. This allows us a deeper insightinto the evolutionary path galaxies have followed since the earliest times. This can beachieved using a range of techniques, including Colour-Magnitude Diagram analysis andspectroscopic abundances. By observing large and small nearby galaxies in detail wecan hope to see evidence of galaxy building processes. Recent signs of this are bursts of

20 chapter 2: Fornax and the Local Group

Figure 2.1: Schematic representation of the Local Group (Grebel 1998).

star formation, tidal debris and on-going mergers, but to find evidence of similar eventsin the distant past we need to uncover more deeply hidden information. One methodis to look for unique chemical signatures in stellar abundance patterns which reveal thedetailed evolutionary history of star formation in galaxies and can be used to determinehow much different galaxies have in common with each other throughout the history ofthe Universe and thus if the assumptions of hierarchical galaxy formation are valid.

Dwarf galaxies offer us the opportunity to study the star formation history and chem-ical evolution of complete systems that are quite different to the MW and likely to bemore similar to (the metal poor and small) galaxies found in the early universe. Theirsmall size also means that to a first approximation, they can be considered as chemicallyhomogeneous “single cell organisms”, creating stars as more of a single unit than a largergalaxy such as the Milky Way; largely unaffected by complexities such as spiral arms,and distinct components such as disk, halo and bulge. There is however the complicationthat it is relatively easy for a small galaxy to loose metals during a supernovae explosions

2.2: Fornax dSph 21

Figure 2.2: From Coleman& Da Costa (2005), thedistribution of Fornax RGBstars, where each star hasbeen convolved with a Gaus-sian of width 10′. The outershell of Fornax is clearly vis-ible, located 1.3◦ north-westof the centre. The first shellis too close to the centre to bevisible.

(e.g. Mac Low & Ferrara 1999). Due to their proximity to the gravitational potentialwell of bigger galaxies (like the Milky Way or M 31 for Local Group Galaxies), dwarfgalaxies are also more likely to loose gas than to attract it and this may explain the dif-ferent characteristics of dSphs and dIrrs (Einasto et al. 1974). When a dwarf galaxy fallswithin the gravitational influence of a larger galaxy it may loose gas, stars and maybeeven globular clusters to the larger host galaxy. It is also plausible that tidal forces willdrive the star formation events (e.g. Mayer et al. 2001). There is clearly a range of dwarfgalaxy properties in the Local Group, and most recently there is evidence that we haveoverlooked a large number of extremely faint dwarf galaxies around the Milky Way (e.g.Belokurov et al. 2006a,b; Willman et al. 2005a,b; Zucker et al. 2006a,b).

2.2 Fornax dSphThe Fornax dSph galaxy is a relatively isolated, dark matter dominated dwarf galaxywith a total mass∗ of 108 − 109 M� (Walker et al. 2006, Battaglia et al. in prep.), at adistance of roughly 135 kpc (Bersier 2000). It is well resolved into individual stars, andcolour-magnitude diagram (CMD) analyses have been made going down to the oldestmain sequence turn-offs (e.g. Stetson et al. 1998; Buonanno et al. 1998; Saviane et al.2000; Gallart et al. 2005). In common with most other dSph, Fornax has no obviousH i associated to it at present, down to a density limit of 4 × 1018cm−2 in the centreand 1019cm−2 at the tidal radius (Young 1999). Unusually for dwarf galaxies, Fornaxcontains five globular clusters (see section 2.3).

∗ luminous mass ' 7× 107, Mateo et al. (1991)


Figure 2.3: From Dinescu et al. (2004) proper motion study, the projection of the orbit(gray line) of Fornax (F). The black lines represents the orbital paths of Fornax and theLMC over the last Gyr. The dashed line represents the Galactic plane. The MagellanicStream is represented with H i column density contours (from Putman et al. 2003), downto a column density of 1019 cm−2. The Sculptor dSph (S) and Phoenix dwarf (P) arealso marked on this plot.

Traditionally, dSphs are considered to be simple, uniform spherical systems. How-ever, in a wide field photometric survey of Fornax, a small overdensity of stars was foundlocated approximately 17′ (or 670 pc) south-east from the centre of Fornax apparentlydominated by a relatively young stellar population with an age of ∼2 Gyr (Coleman et al.2004) It is possible that this might be a shell structure, something previously unseen in adwarf galaxy, which may be the remnant of a merger with a small, gas-rich system thatoccurred approximately 2 Gyr ago, although a detailed study by Olszewski et al. (2006)suggests that the metallicity of this stellar population is the same as Fornax. A secondlarge shell-like structure has been discovered 1.3◦ north-west from the centre of Fornax,outside the nominal tidal radius (Coleman & Da Costa 2005, see Figure 2.2).

Clearly Fornax exists in a complex environment. Proper motion studies of Fornax, us-ing a combination of photographic plate material and HST Wide Field Planetary Camera2 data suggest that Fornax crossed the Magellanic plane ∼190 Myr ago (Dinescu et al.2004, see Figure 2.3). This crossing appears to roughly coincide with the terminationof all star formation in Fornax (Stetson et al. 1998). It is possible that ram pressurestripping of the ISM may have caused the end of star formation in Fornax. There areH i clouds found all along the proposed orbit of Fornax consistent with stripped mate-rial from Fornax as it crossed the orbit of the Magellanic Clouds (Dinescu et al. 2004).However, there remains a distance discrepancy between Fornax (135 kpc) and the LMC(50 kpc) which makes it difficult to be sure if there was any interaction at all.

2.2: Fornax dSph 23

Figure 2.4: (top): RelativeSFR for Fornax, taken fromTolstoy et al. (2001). (bot-tom): SFR of Fornax innerfield, taken from Gallart et al.(2005).

Using CMDs from several sources, Tolstoy et al. (2001) constructed a schematic starformation history for Fornax dSph. This is presented in the top panel of Figure 2.4.More recently, a new star formation history has been published by Gallart et al. (2005),for the inner field of Fornax, which we reproduced in the bottom panel of Figure 2.4.The two plots do not agree very well, although both show a peak in star formation at∼4 Gyr. The relative number of young and ancient stars differs significantly. This mightbe due to studies covering different regions of Fornax which we now know has quite a lotof spatial variation in stellar population (Battaglia et al. 2006). These differences requirefurther study.

Low resolution spectroscopic studies of individual stars have been made to determineCa II triplet metallicities for samples of ∼30 stars (Tolstoy et al. 2001), ∼100 stars (Pontet al. 2004) and most recently ∼600 stars (Battaglia et al. 2006). These studies haveshown that Fornax contains a relatively metal-rich stellar population and has a complexstar formation history where the majority of stars have been created at intermediate ages2− 6 Gyr ago with a peak at 5.4±1.7 Gyr ago. (Saviane et al. 2000). Fornax also has ayoung stellar population (<1 Gyr) as well as an ancient one (10-12 Gyr) and stars witha range in metallicity going from −2.8 dex to solar.

Detailed abundance analyses based on UVES high resolution spectroscopy have alsobeen carried out (Shetrone et al. 2003; Tolstoy et al. 2003) but these were limited tothree stars. In chapter 5, we present our abundance results for nine stars belonging tothree of the Fornax globular clusters, and in chapter 6, the abundances of an additional81 stars belonging to the central (25′) field of Fornax.


Figure 2.5: The HST Colour Magnitude Diagrams of three of the Fornax GlobularClusters, from Buonanno et al. (1998). The RGB stars for which we obtain HR spectra(see chapter 5) are marked.

2.3 Globular Clusters in FornaxGlobular clusters are rarely found associated to low mass dwarf galaxies. It is not clear ifthis is because low mass galaxies loose their globular cluster populations, or if they typ-ically don’t form stars actively enough to warrant a globular cluster population. Fornaxand Saggitarius are the closest dwarf galaxies with globular clusters. Contrary to theSagittarius dSph which is obscured by dust, and confused by merging with our Galaxy,the Fornax dSph is high above the Galactic plane and offers a uniquely useful target forinvestigation. The next example of a dwarf galaxy with globular clusters is WLM (Wolf-Lundmark-Melotte), which is nearly 1 Mpc away, with only one globular cluster. Bystudying the Fornax globular cluster population we can determine if all globular clustersshare the same properties regardless of the environment in which they were created orif there are differences related to the size, type or location of their host galaxy. At thevery least globular clusters are a part of the overall picture of the star formation historyof a galaxy.

Unlike most other dSphs, Fornax contains five globular clusters (Shapley 1938; Hodge1961), .1 kpc from its centre. This highly unusual specific frequency (∼70) is an orderof magnitude higher than the expected value for a galaxy of its size (. 5, Harris 1991).According to Goerdt et al. (2006), in a cuspy cold dark matter halo, Fornax GCs shouldsink to the centre within a few Gyr, raising the question of how these old GCs could havesurvived to the present epoch. As Goerdt et al. (2006) show, a solution to this timingproblem is to adopt a cored dark matter halo. Under these conditions, it will take the GCsmany Hubble times to sink to the centre, as they will stall at the dark matter core radius.

Globular clusters are typically associated with the oldest stellar population compo-nent of a galaxy. We do not know for sure under which conditions they form and survivebut they are generally assumed to be a ubiquitous old population associated with theepoch of galaxy formation. Every large galaxy (spiral or elliptical) appears to have apopulation of very old globular clusters (Harris 1991). It is commonly believed thatglobular clusters are formed during periods of exceptional star formation, such as during

2.3: Globular Clusters in Fornax 25

the initial formation of the galaxy (Searle & Zinn 1978) or during a major merger (whichmight be the same thing). However, this view breaks down in the Magellanic Cloudswhere there are (at least) two distinct populations of globular clusters (one young andrelatively metal-rich, and the other old and metal-poor) and there is no similar signaturein the field star stellar population (van den Bergh 1981).

The Fornax dSph GCs, similarly to Galactic GCs, are single-age stellar populations.Their globular-cluster-like ages have been determined by isochrones fitting on HST CMDs(see Figure 2.5) going down to oldest main sequence turn-offs and are the same to within± 1 Gyr, with the possible exception of cluster 4, which is buried in the centre. Somestudies have found it to be younger by about 3 Gyr (Buonanno et al. 1998). The metallic-ities of the clusters vary, as summarised in Strader et al. (2003), but are more metal-poorthan the average for the field stellar population (by a factor of more than ∼1 dex), witha bluer RGB, well populated blue horizontal branches (HB) and a range of HB morphol-ogy (Buonanno et al. 1998, 1999). Thus the Fornax globular clusters represent quite adifferent stellar population to the Fornax field stars (Stetson et al. 1998; Buonanno et al.1999; Saviane et al. 2000).

Chapter 3Using stellar atmospheric models todetermine chemical abundances

In order to quantify the different chemical elements present in a stellar atmosphere, weneed to describe the star in a physical way. The light that we receive from a star

comes through it’s atmosphere, more specifically the photosphere. Photons emerge fromthese transparent layers of gas, releasing the energy produced by the thermonuclear re-actions in the star’s opaque centre. The temperature, pressure and chemical compositionof the atmosphere will determine the features of the star’s spectrum. Absorption linesare created when a particle (atom or molecule) absorbs a photon from the emerging fluxat a specific wavelength. Each different chemical element will absorb photons at specificwavelengths and by measuring the relative depth of these absorption lines we can deter-mine the abundance of that particular element.

The following paragraphs are not meant as a summary of the physics of radiativetransport in stellar atmospheres. The theory of stellar atmospheres is a well developedpart of astrophysics and the detailed physical processes are well described in standardtextbooks such as Gray (1992, chapters 5–14) or Carroll & Ostlie (1996, chapters 9–10). On the other hand, some general background will help to understand the analysistechniques used in this thesis. Sections 3.1 – 3.4 therefore give an overview of the mostimportant concepts and terminology that will be frequently used in the subsequent chap-ters. In order to facilitate reading, Table 3.1 lists the physical constants used in thischapter.

3.1 Describing the stellar atmosphereStellar atmospheres are low density gas, so the “ideal gas law” may be used to relate thepressure, density and temperature. Here I summarise the physical description of the fluxemerging from an ideal atmosphere and the different absorption sources present in suchan atmosphere.

28 chapter 3: Using stellar atmospheric models ... chemical abundances

Figure 3.1: Volume element illustrating flux and intensity in a stellar atmosphere.Adapted from Gray (1992), figure 5.1 and 5.3.

Table 3.1: Constants used in this chapterName symbol value unitsSpeed of light c 2.99792458× 108 m s−1

Planck’s constant h 6.63× 10−34 J sBoltzmann’s constant k 1.38× 10−23 J K−1

8.62× 10−5 eV K−1

Stefan-Boltzmann’s constant σ 5.6705× 108 W/m2 K4

Gravitational constant G 6.672× 10−11 m3 kg−1 s−2

Mass of the electron me 9.11× 10−31 kg

3.1.1 The flux

Lets consider a cylindrical volume element of surface dA and thickness dx, as shown inFigure 3.1, radiating at a frequency ν and intensity Iν . The radiation is emitted in adirection θ with respect to the cylindrical axis, per unit area, unit solid angle (dω), unittime and unit frequency. The basic equation that describes radiative transfer in a caselike this is the following:

dIν

dτν= −Iν + Sν (3.1)

where Sν is the source function (Sν = jν/κν), jν and κν are the emission and absorptioncoefficients, τν is the line of sight optical depth (τν =

∫κνρ dx, with ρ representing the

density of matter in the unit volume.) So the flux Fν that is crossing the volume elementper unit time and frequency is defined by:

3.1: Describing the stellar atmosphere 29

Fν =∮

Iν cos θ dω (3.2)

Although the basic radiative transfer relation (3.1) looks extremely simple, this simplicityis very delusive, mainly because the quantity κν involves a large amount of complexphysics. In order to solve the transfer equation and arrive at the atmosphere structureand the photospheric spectrum of the star a number of simplifying assumptions areneeded:

Hydrostatic Equilibrium:

This is the case when pressure forces balance gravity. There is no expansion and nosignificative mass loss.

Thin atmosphere:

The thickness of the photosphere is small compared to the radius of the star. Thus, weneed only consider the atmosphere as a superposition of parallel planes or “onion shells”(layers) with a single (1D, radial) dimension describing the structure. We may thusassume that the variation of gravity over the thickness of the photosphere is negligibleand we can approximate the gravity as a constant.

Local Thermodynamic Equilibrium (LTE):

We assume that LTE is a valid approximation for each volume element in the atmosphere.Every layer has a unique temperature (T = T (τν)) and the source function is the Planckfunction:

Sν = Bν(T ) =2hν3

c2

1exp(hν/kT )− 1

(3.3)

where c is the speed of light, h is the Planck constant and k is the Boltzmann constant.The LTE approximation allows us to use the following two laws:

• Boltzmann’s Law: To know whether a particular line may occur, you have toknow the relative populations of the excited states of the particles in the gas. Therelative population of excited states in a gas in thermodynamic equilibrium is givenby the Boltzmann Excitation Distribution. The number of atoms of energy level nper unit volume Nn is proportional to the total number of atoms (N) of the samespecies:

Nn

N=

gn

Un(T )exp

(−χn

kT

)(3.4)

where gn is the statistical weight of the nth level, χn is the excitation potential of thenth level and Un(T ) is the partition function of the particle in a gas of temperatureT and is defined as: Un(T ) = Σgi exp(−χi/kT ). It is often the case that χn isexpressed in eV and the term (1/kT ) is often expressed as θ = log e/kT = 5040/Twhich lead to:


Nn

N=

gn

Un(T )10−θχn (3.5)

• Saha’s Law: In order to describe an absorption line, we need to know what fractionof the atoms of a particular element are in the ionization state corresponding tothe line. Saha’s law describes the distribution of particles of the same species indifferent ionization states. The ratio of atoms in ionization state i and i + 1 isrelated to the electronic pressure (Pe) and temperature T of the gas :

Ni+1

NiPe =

(2πme)3/2(kT )5/2

h3

Ui+1(T )Ui(T )

exp(− χi

kT

)(3.6)

where me is the mass of the electron and χi is the ionization potential of the ion inthe state i. The Pe term in that equation explains why stellar spectra are sensitiveto pressure. The assumption of LTE is a very important simplification of the gen-eral problem, as it allows us to calculate the source function, the population of theatomic energy levels and the ionization equilibria from only a small number of freephysical parameters. In very thin extended atmospheres, or in the case of strongabsorption lines which are formed in high atmospheric layers, the LTE assumptionbreaks down. The calculation of the excitation and ionization equilibria then be-comes enormously more complicated because all interactions between matter andradiation have to be considered in detail.

Radiative Equilibrium

In the top layers of any stellar atmosphere, all the energy is carried by radiation. Conser-vation of energy tells us that the energy absorbed by one layer in the atmosphere mustbe re-emitted to the next, or in other words, the flux must be constant (F(x) = F0)throughout the atmosphere. In the case of a 1D model, we have:

ddxF(x) = 0 (3.7)

where F(x) is the total flux (in W/m2). When all the energy is carried via radiation, wehave: ∫ ∞

0

Fνdν = F0 = constant = σT 4eff (3.8)

where σ is the Stefan-Boltzmann constant and Teff is the black body temperature of thestellar atmosphere.

3.1.2 The absorption coefficientAny process that captures or prevents photons from being emitted by the atmospherewill contribute to the absorption coefficient (or opacity). This includes scattering as wellas absorption of photons by atomic electrons making level transitions. The absorptioncoefficient (κν) of a gas is obviously going to be frequency dependent.


Continuous absorption

This is the sum of the absorption resulting from many physical processes. The wave-length dependence of the continuous absorption coefficient shapes the continuous spec-trum emitted by a star. Photoionization, when a photon has enough energy to ionise anatom (bound-free absorption) is a source of continuous opacity. Also free-free absorption(when a free electron in the vicinity of an ion absorb a photon) contributes to the contin-uous opacity of the star. Electron scattering (Thompson, Compton, Rayleigh) can alsodivert photons from an incident light source, so they also contribute to the continuousabsorption. Hydrogen, being the most abundant element, is also the main contributorto the absorption coefficient. In cool stars like those of our sample (∼ 4000 K) most ofthe continuous absorption in the visible and infrared part of the spectrum is due to thenegative hydrogen ions H− (hydrogen atoms with one very loosely bound extra electron),while “metals” start to dominate the UV part of the spectrum.

Specific absorption

Absorption specific to the line, (bound-bound transitions) occurs when an electron inan atom or an ion makes a transition (by absorbing a photon) from one orbital to an-other. It is, by definition, very wavelength specific, corresponding to the energy of thephoton that was absorbed. The depth and width of this absorption line is related tothe transition probability, the population of the lower energy level, and the abundanceof the element that absorbed the photon, but also to some intrinsic effects not relatedto the abundance. Natural broadening is caused by Heisenberg’s uncertainty principle,where the orbital energy cannot have a precise value, allowing for photons of slightlydifferent wavelength to be absorbed. This results in a non-discrete (fuzzy) energy level.Thermal (or Doppler) broadening is caused by the fact that atoms are in thermal motion,producing a range of line of sight velocities. This motion will change the observed fre-quencies (Doppler shifting) of the absorbed photons, making the line broader. There isalso pressure broadening, caused by the electric field of a large number of (close by) ionsand collisional broadening, when the orbitals of an atom are perturbed due to collisionwith a neutral atom.

Macro and micro turbulence are two broadening mechanisms that act on scales thatare large (macro) or small (micro) compared to the mean free path of the photons.Microturbulence can be considered as an additional thermal velocity. When the line ofsight goes through many cells of motion (turbulence cells) in the photosphere, the velocityof the cells will modify the line profile in the same way as the particle distribution. Itis approximated to be isotropic (Gaussian) and can be included directly into the lineabsorption coefficient with a convolution, as detailed in chapter 18 (p.405) of Gray (1992).There is macroturbulence when the turbulence cells in the photosphere are large enoughso that a photon will stay in the same cell from the time it is created to the time itleaves the star. Each of these cells will have the same Doppler shift, corresponding to thevelocity of the cell, therefore acting in a way similar to rotation, which can be applied as aconvolution of the emergent spectrum by an appropriate function (Gaussian or other). Tosummarise, micro turbulence acts on the absorption line profile, like a thermal componentdesaturating strong lines while macro turbulence acts on both strong and weak lines inthe same way by smearing them out over a frequency range.


Figure 3.2: Electronic pressure (top), gas pressure (middle) and temperature (bottom)as a function of optical depth. This figure present models with constant log g and [Fe/H]in order to illustrate the Teff dependence of the models. Teff start at 3800 K (solid line)and increase by 100 K each time to reach 4200 K (dotted line). The models used arethose of MARCS 2005 and Plez 2005, presented in section 3.3.2.


Figure 3.3: Electronic pressure (top), gas pressure (middle) and temperature (bottom)as a function of optical depth. This figure present models with constant Teff and log g inorder to illustrate the [Fe/H] dependence of the models. [Fe/H] start at -2.5 dex (solidline) and increase by 0.5 dex each time to reach -0.5 dex (dotted line). The models usedare those of MARCS 2005, presented in section 3.3.2.


Figure 3.4: Electronic pressure (top), gas pressure (middle) and temperature (bottom)as a function of optical depth. This figure present models with constant Teff and [Fe/H]in order to illustrate the log g dependence of the models. The log g start at 0.0 dex (solidline) and increase by 0.3 dex each time to reach 1.2 dex (dotted line). The models usedare those of MARCS 2005, presented in section 3.3.2.

3.2: Determining Stellar Atmospheric parameters 35

3.1.3 Stellar atmospheric modelsStellar atmospheric models are a tabulation of physical parameters used to represent theconditions inside an atmosphere. Models are typically given as the electronic pressure(Pe), the gas pressure (Pg), the temperature (T ) and the optical depth for photonswith λ=5000Å (τ5000) for several layers (∼50) of a stellar atmosphere. This is shown inFigures 3.2, 3.3 and 3.4 where we plot Pe (top), Pg (middle) and T (bottom) as a functionof τ5000 for five different set of parameters, varying Teff , [Fe/H] and log g respectively,sampling the full range of stellar parameters we used in chapter 6. It is customary instellar atmosphere work to use log g as equivalent for the pressure in the atmosphere(which can be done if hydrostatic equilibrium is valid). In these three figures, wherewe can see the similarity in the shape of the curves when only one parameter changes,allowing us to interpolate between two curves to get the exact parameter needed for ourmodel. As can be seen from the plots of T versus τ5000, the parameter that will mostinfluence the line formation is Teff . Especially in the region where most of the lines areforming, (−1 < τ5000 < 1), a change of 100 K will change the T (τ) relation much morethan a change in log g and/or [Fe/H]. Therefore it is critical to have stellar models madefor the Teff corresponding to the star observed in order to produce accurate abundances.These models are needed as an input for the line formation code used to derive theabundance, as describe in section 3.3.3.

3.2 Determining Stellar Atmospheric parametersIn the previous section, we have shown that we can simplify our stellar atmosphere modelso that it can be described using only a few parameters. We will describe them (and howto derive them) in this section.

3.2.1 Effective Temperature (Teff)The effective temperature, Teff , is the temperature of a black body radiating as the star,F = σ Teff

4. There is more than one way to determine the Teff of a star, and I willdescribe the two methods we used.

Photometric colour

A powerful method to obtain the Teff of a star, is the InfraRed Flux Method, (IRFM)described by Blackwell & Lynas-Gray (1998). It provides an accurate procedure to derivestellar angular diameters and effective temperatures by measuring the monochromaticflux at an infrared frequency and the bolometric flux. It then uses theoretical atmosphericmodels to estimate the monochromatic flux at the star’s surface (the infrared flux hasa small dependency on the Teff). The IRFM is an iterative procedure: from a firstguess of the Teff , the angular diameter is deduced and used to derive an improved Teff .Practically, the IRFM is not easy to use since measuring a stellar angular diameter isnot always possible. An empirical method, calibrated on the IRFM, has been developedto give a relation between photometric colours (like V − I, V −K) and Teff . The generalmethod and correction polynomials are described in the series of papers by Ramírez &Meléndez (2005), Alonso et al. (1999a,b, 2001), and references therein, and we will usethese calibrations in the following chapters.


Excitation equilibrium

We define Teff such that the abundance of an element is independent of the excitationpotential (χex) of the individual lines. Obviously, in principle, all the lines of an elementshould give the same abundance for a given star. In practice, there is a (small) scatteraround an average value. A Teff which is incorrect will affect the weak excitation poten-tials more than the strong potentials. It will also change the gradient of the temperatureand the optical depth (T and τ5000) and since lines with different χex are not all formingat the same depth, the resulting abundance will be different. If a correlation betweenabundance and excitation potentials occurs, it is a sure sign that Teff has been incorrectlydetermined for the star. In order to use this method, we need many lines of a single el-ement sampling a range of χex. Figure 3.5 (middle panel) illustrates this for a samplestar of our FLAMES dataset, with the proper Teff , where the slope is effectively zero.Contrasting with this figure, we show in Figure 3.6 the same figure but with a Teff thatis 400 K higher than our ”correct“ one. The precision with which Teff can be determineddepends upon the resolution, the choice and number of lines and signal to noise of eachspectrum used. Usually, we use ≈50 Fe i lines to determine the Teff of a star.

3.2.2 Surface Gravity (log g)The surface gravity of a star of mass M? and radius R? is define as g? = GM?/R2

?, whereG is the gravitational constant. In solar values, we get g? = g�(M?/M�/)/(R?/R�)2.There exists several methods (isochrones, pressure broadening in the wings of stronglines) to estimate the gravity of a star (which we often use in log scale, log g) and in thissection we describe the two different methods we used to calculate it.

Photometric

We can use photometry to estimate the surface gravity of a star if we know the mass,the Bolometric magnitude (distance modulus + a bolometric correction) and the Teff

(two or more photometric colours). We do so using the relation linking the luminosity,radius and temperature of a star (L = 4πR2σT 4) and Bolometric magnitude definition(MBol? −MBol� = −2.5 log(L?/L�)) we get:

log g? = log g� + logM?

M�+ 4× log

Teff?

Teff�+ 0.4× (MBol? −MBol�) (3.9)

Spectroscopic (Ionization Equilibrium)

For stellar types F, G or K, we can easily measure elements in two ionization states,like Fe i and Fe ii or Ti i and Ti ii. For a given star and a given element, there shouldbe a single value for the abundance, no matter if the abundance is determined from theneutral or the ionized state. We can then iterate on the gravity of our model until theabundance of Fe i and Fe ii are the same, constraining the stellar gravity. By definition,gravity is related to the gas pressure (Pg ∝ g2/3) and therefore to the electronic pressure(Pe ∝ g1/3) since Pg ∝ P 2

e . From the Saha’s equation (3.6) and the cool stars case (ourcase), where the number of atoms of Fe i � Fe ii, we can state that Fe i, the dominantspecies, will depend on 1/Pe and Fe ii, the minority species, with the majority of atomsin the state i−1 = 1, will depend on 1/P 2

e . This is what makes the ionization equilibrium

3.2: Determining Stellar Atmospheric parameters 37

Figure 3.5: Fitting a model (MARCS 2005) to a star in our FLAMES sample (BL239),showing the [Fe i/H] (full symbols) and [Fe ii/H] abundances (empty symbols) as a func-tion of λ, (top) χex (middle) and EW (bottom). The text above the top plot gives theparameters of the model used: star name, Teff (T), log g (g), vt (v), metallicity (z), theaverage [Fe i/H] (also plotted with a dashed line) and the number of Fe i lines used. Thethick line in the middle and bottom plots are linear regressions, with their respectivecoefficients and associated errors on top of each plot. The different symbols for the Fe ilines are related to their equivalent width, as shown in the bottom panel.


Figure 3.6: Same as Figure 3.5 but with a Teff that is 400 K higher than the ”correct“one. The slope in the middle plot, although only significative at the 1.8σ level, shows thatthere is a tendency for weak χex to produce higher abundances. This extreme change inTeff changes the derived [Fe i/H] by ∼0.2 dex, and shift the [Fe ii/H] abundances (emptysquares) from relatively comparable to the [Fe i/H] to way below them, another hint thatthis temperature is not appropriate for this star.

3.3: The abundance determination 39

a good tool to constrain gravity. This method presumes that non-LTE effects are notmodifying the ionization equilibrium. This is an assumption that is not always correct,especially at low surface gravities or metallicities, where Fe is known to be overionisedin non-LTE (Asplund 2005). This translates into an underpopulation of Fe i levels withrespect to what was predicted by LTE and therefore an [Fe i/H] abundance lower than[Fe ii/H]. Although it is claimed by many authors (including Asplund (2005)) that Fe iiis immune to departures from LTE in late-type stars, the typical number of Fe i linesobserved versus Fe ii lines (factor ∼10) makes Fe i a more reliable measure of Fe thanFe ii.

3.2.3 MetallicityEach stellar atmosphere model is computed with a parameter representing the chemicalcomposition of the star. It is often referred to as [Fe/H], although it does not onlyrepresent the contribution of Fe atoms. It represents the electronic pressure (Pe) in theatmosphere of a star with the same chemical element ratios as the Sun, scaled to a given[Fe/H]. It is the abundance of the elements that contribute to the continuous absorptionproperties of the atmosphere. A higher metallicity will increase Pe in the atmosphere bycontributing extra electrons. Some models have different [α/Fe] ratios to represent starsthat are systematically different from the sun.

3.2.4 Microturbulence velocityThe microturbulence velocity, vt affects the lines by broadening and hence desaturatingthem. It is caused by small cells of motions in the photosphere and is treated like anadditional thermal velocity in the line absorption coefficient. The desaturation effectdepends on the strength of the line, where only strong lines are affected. Weak lineswill not be affected by desaturation since increasing the vt will broaden the line andmake it shallower, conserving the equivalent width. In this regime, the abundance isproportional to the EW . But for a saturated line, increasing the vt will widen thewavelength range covered by the absorption, thus desaturating the line: the equivalentwidth is not conserved anymore. This is detailed in in chapter 18 of Gray (1992). Typicalvalues for the vt are 1-2 km/s for low-mass giants. We can determine vt for a star bymaking sure that for a single element, the abundance is independent of the EW ofthe line, as illustrated in Figure 3.5 (bottom panel) for a sample star of our FLAMESdataset, for which the vt has been chosen correctly for the observed spectra (negligibleslope). Again, Fe i, having many observed lines, is the most suitable element for thisdetermination.

3.3 The abundance determinationAfter describing the physics of a stellar atmosphere and parameterising it, we need totransform the absorption lines into chemical abundances. We will first need to mea-sure the equivalent widths of the absorption lines, and then use a stellar atmospheremodel with the right parameters for each individual star before being able to derive anabundance.


Figure 3.7:The equivalentwidth of anabsorption lineis defined asthe width of arectangle thathas an areaequal to theline, as illus-trated in grey.Adapted fromFigure 9.18 ofCarroll & Ostlie(1996)

3.3.1 Measuring the equivalent widthsThe first step in determining the abundance is the actual measurement of the strengthof each absorption line. We refer to this value as the equivalent width (EW , in text orW , in equations), it corresponds to the total absorption coming from a line, and it isdefined in the following way:

EW = W =∫ +∞

−∞

Fc − Fν

FcdFν (3.10)

Where EW is the width of a rectangle of depth 100% (going from 0 to 1) in a normalisedspectrum that covers the same area as the real line. This is illustrated in Figure 3.7,where Fc is the flux level of the continuum (normalised at 1), Fλ = Fν is the flux at thefrequency ν = c/λ, with c = the speed of light.

3.3.2 The Stellar Models usedIn 2005, a major improvement in the models available occurred with the release of newMARCS spherical stellar models∗ which are described in in Gustafsson et al. (2003). Forchapter 5, (which was made in prior to 2005) we used models from Plez (2000, 2002).For chapter 6, we used the new MARCS 2005 models extended by Plez (2005) to coverthe range of stellar parameters of our sample. Here is a summary of the models used:

Plez 2000-2002

• Geometry: Plane-parallel approximation

• Temperature: 3800 ≤ Teff ≤ 5200 K in steps of 200 K

• Gravity: 0.5 ≤ log g ≤ 4.5 dex in steps of 0.5 dex∗ http://marcs.astro.uu.se/

http://marcs.astro.uu.se/

3.3: The abundance determination 41

• Metallicity: −4.0 ≤ [Fe/H] ≤ −1.00 dex in steps of 0.25 dex

• Alpha: Enhanced, [α/Fe] = 0.4

MARCS 2005

• Geometry: Spherical


• Gravity: 0.0 ≤ log g ≤ 3.5 dex in steps of 0.5 dex

• Metallicity: −1.5 ≤ [Fe/H] ≤ +1.00 dex in steps of 0.25 dex

• Alpha: Standard, [α/Fe] = 0 at [Fe/H] = 0, +0.1 for each -0.25 dex until it reaches+0.4 at [Fe/H] ≤ -1.0.

Plez 2005

• Geometry: Spherical


• Gravity: same as MARCS 2005

• Metallicity: −3.0 ≤ [Fe/H] ≤ −1.50 dex in steps of 0.5 dex

• Alpha: Poor, [α/Fe] = 0.00 for all models.

Models are interpolated for all parameters, including [α/Fe] when mixing standard mod-els with α-poor ones, in order to create the correct model for individual stars. We usedmodels with two types of geometry: plane-parallel and spherical, but for a given sample,we used either one or the other for the entire analysis. The geometry affects the abun-dance in two fundamental ways, namely the line formation (discussed in section 3.3.3)and the model atmosphere structure. Spherical geometry in the model structure is abetter representation of reality but as they are relatively new models they have notbeen used intensively in the literature, since prior to the MARCS 2005 models, the mostcommonly used models for abundance analysis in giants stars were those of Gustafssonet al. (1975), in plane-parallel. This makes a direct comparison with previous work morecomplex but since these models became available, we decided to start to use them. Anoverview of the difference between models with spherical geometry and plane-parallelapproximation is presented in Heiter & Eriksson (2006), where they compare the effectof using the plane-parallel approximation in the model atmosphere structure and/or theline formation code (p_p and s_p) using fully consistant spherical geometry (s_s) in theabundance analysis. They show that lines with different χex will not behave in the sameway (introducing a bias on the determined Teff), lines with different EW will also showa different behaviour (affecting the vt) and that lines of different ionization state (Fe iversus Fe ii) will also react differently to a change in geometry (bias in log g). They givethe maximum combined systematic error caused by different geometry in the s_p casewith respect to the s_s case to be of the order of -0.1 dex, while for the p_p case, thedifferences are up to +0.35 dex. Thus using spherical models can make a big differencein the abundance determination, much more than the way we treat the line formation(which in our case is plane-parallel).


3.3.3 Computing the abundancesAs explained in detail in chapter 14 of Gray (1992), for weak lines (dominated by Dopplerbroadening) we can show that:

log(

Wλ

λ

)= log

(πe2

mec2

Ni/N

U(T )NH

)+ log A + log(gf λ)− 5040

Tχ− log(κν) (3.11)

where Wλ is the equivalent width of the line, e is the charge of the electron, e = −1.60×10−19C; Ni/N is the ratio of the number of atoms of a particular element in the ionizationstate i with respect to the total number of atoms of that element, NH is the number ofhydrogen atoms per unit volume, A = N/NH is the abundance of the specific elementrelative to hydrogen, Un(T ) is the partition function, defined in equation 3.4, κν is thecontinuous absorption coefficient and gf is the transition probability∗. Note that thefirst term on right hand side of the equation is constant for a given star and a given ion.This equation gives us some general information about the abundance of an element ina star:

• for weak lines, the equivalent width (W ) varies in a linear way with abundance.

• the abundance of an element (A) varies with the inverse of the temperature (5040/T ).

• the abundance depends linearly on the gf -values.

The line formation code

To calculate the abundances we use CALRAI, developed by Spite (1967) with manyimprovements over the years, which uses the plane-parallel approximation. Once wehave determined the appropriate stellar parameters for a given star, we can interpolatethe model from our grid and use the measured EW s to determine the abundance ofthe different elements. This is an iterative process, where the software will vary theabundance of a given element until it is consistent with the observed EW s. Once thishas been done for all the lines of a given element, we obtain a distribution of abundancesfor each element. The more lines of a single element we have the greater the reliability ofthe abundance derived. The different lines do not necessarily behave in the same way withrespect to abundance. Weak and moderately strong lines will vary in an almost linearway with respect to the abundance. When lines become saturated (without prominantwings), a strong variation in abundance will be almost insensitive to the EW . A curveof growth can be used to illustrated this dependence.

The Curve of Growth

For a given abundance (α?) of a specific element, the EW s (or Wλ) of a given line willvary as a function of the gf -value (the transition probability, eq. 3.11). When we have aweak line, log (Wλ/λ) will vary linearly with log (α?gf). We can define Γ? for weak lineswhen this equation is true:∗ where g is the statistical weight (2J+1, J is the inner quantum number) of the lower level, and f is

the oscillator strength

3.4: The line list 43

log(

Wλ

λ

)= log (α?gf) + log Γ? (3.12)

We can calculate Γ? for each line as a function of the model used, the element and itsionization state and its χex. A curve of growth is a plot of log (Wλ/λ) versus log (α?gf)+log Γ?. Figure 5.3 of chapter 5 illustrates this for one star in our sample. Since we don’tknow a priori the value of α?, we use the solar value, α� and the equation becomeslog (α�gf) + log Γ?. If we consider Fe for example, then all the lines of Fe i and Fe iiwill be aligned on a curved shifted by log (α?)− log (α�) with respect to the theoreticalcurve of growth. The value of this shift is, by definition, [Fe/H]. When we compare theobserved curves of growth for Fe i and Fe ii (one for each state) we can verify if the chosengravity is representative of the star. The gravity of the model will be representative ofthe star if the two different (ionization) states fall on their respective curve. By plottingdifferent symbols for lines of strong and weak χex, we can confirm our choice of Teff .Lines of different χex should be randomly distributed higher and lower than the curve.If it’s not the case, it is a sign that the Teff chosen for the model is not a good match tothe star (so we have to change it). Also, if the vt is not the right one for the star, thenon-linear part of the curve of growth will not fit the observed lines.

3.4 The line listHigh resolution spectroscopy can provide accurate measurements of numerous absorptionlines for many different chemical elements and hence is an accurate method of determiningdetailed abundance patterns in a star. The line list is a critical part of the analysis, andbuilding a proper line list is a complex task. Lines needs to be chosen carefully, makingsure that they have a reliable gf -values and are sufficiently isolated from their neighboursat the resolution of the observations to be accurately measured. A line that is isolatedfor a star of a given metallicity (or Teff) can be too blended to be useful at anothermetallicity. Also, a line that is isolated in a high resolution instrument (like UVES forexample) can be blended in a lower resolution one (like GIRAFFE). In addition, becauseof different wavelength coverage, some lines are available for one instrument but not foranother. So a line list needs to be adapted for the data set it will be used on.

3.4.1 Building a line listOne way to build a line list is to start from basic reference line list and add/substractlines from it. Such a basic line list can be found in recently published work made on sim-ilar objects. Then, extra lines can be added from other work and/or web-based atomiclibraries like NIST∗, VALD† or Kurucz‡. It is always important to check that there iscompatibility in the average values of the abundances derived from each list. A smallbut noticeable offset in the average value of [Fe i/H] between two lists will add an ar-tificial scatter to the abundances values. The consistency (precision and uniformity) ofthe atomic data used for each list – especially the gf -values – has to be checked before∗ http://physics.nist.gov/PhysRefData/ASD/lines_form.html† http://ams.astro.univie.ac.at/vald/‡ http://cfa-www.harvard.edu/amdata/ampdata/kurucz23/sekur.html

http://physics.nist.gov/PhysRefData/ASD/lines_form.html

http://ams.astro.univie.ac.at/vald/

http://cfa-www.harvard.edu/amdata/ampdata/kurucz23/sekur.html


Figure 3.8: ([Fe i/H]) abundance determination for Arcturus using four different linelists, represented by different symbols. Also plotted are the average abundance for eachof the sets of lines. Lines tagged Shet2003 (filled circles) are from Shetrone et al. (2003),Hill2000 (empty squares) from Hill et al. (2000), Grat2003 (gray diamonds) from Grattonet al. (2003) and Zoca2004 (empty triangles) from Zoccali et al. (2004).

we make a big list out of smaller ones. For some elements, the gf -values are not knownprecisely and the uncertainties are much bigger than what we would hope to achieve withour measurements. The gf -values affect the abundance in a direct way, an error of 0.1on the log gf will affect the abundance by 0.1.

Our strategy was to start with the reference line list from Shetrone et al. (2003),which was optimised for use with UVES on metal poor stars (−3.0 ≤ [Fe/H] ≤ −1.5).For our work on the globular clusters of Fornax, (chapter 5), we used it without anymodifications, as it was already appropriate for our metallicity range and instrumentalresolution. But for the Fornax field stars (chapter 6), because of the higher metallicity ofthe stars and the different (smaller) wavelength coverage, we had to significantly adaptthe line list. We selected lines from Gratton et al. (2003), Hill et al. (2000) and Zoccaliet al. (2004) as potential candidates to be included in our master line list. We used ahigh resolution, (R ≈ 120 000) high signal-to-noise spectrum of Arcturus, downloadedfrom the ESO UVES public archive∗. Arcturus has Teff 4250 K and [Fe/H] ∼ −0.5 and∗ http://archive.eso.org/wdb/wdb/eso/uves/form

http://archive.eso.org/wdb/wdb/eso/uves/form

3.4: The line list 45

is thus a good “template” for our most metal rich stars. On average, as you increase inmetallicity, the strength of a line will increase, making weak lines stronger and saturatingalready strong lines. Our reference line list was optimised (i.e. having both weak andstrong lines) for metal poor stars. In order to optimise our line list also for the richestpart of our sample, we need to add lines that are weak in metal rich stars.

From the Arcturus spectrum we calculated abundances using the four line lists forall the elements and compared the results before accepting new lines. We illustrate thisfor the Fe i lines in Figure 3.8. From this we can determine if the inclusion of a line listwill improve the abundance determination or if it will just add more scatter. If there isnot a big difference in the average value between two lists and/or if a list has a scattercomparable to other lists, we can safely add the new lines to our master list. FromFigure 3.8, we concluded that the lines used by Hill et al. (2000), apart from being toostrong, would only add scatter to our Fe i distribution. Lines from Zoccali et al. (2004)have an average Fe i abundance difference of ≈ 0.2 dex compared to the Shetrone list,a sign that the gf -values of the two lists are not compatible. It is possible to shift thegf -values so that the average abundance comes out the same as the Shetrone one butsince the Zoccali list has almost no weak lines, we didn’t do it and didn’t use their lines.Only the lines from Gratton et al. (2003), because of their relative weakness (30-70 mÅ)and small scatter, were added to our master list. Since the difference in average log gfis small (0.04 dex), we decided to ignore this difference and use the gf -values withoutmodification.

3.4.2 The line by line selectionTo facilitate the addition of “clean” lines to our line list, we used two synthetic spec-tra that have been made with models that have the stellar parameters of Arcturus,Teff=4250, [Fe/H]=0.5, log g= 1.5, one containing most of the known atomic and molec-ular lines convolved to our instrumental resolution and the other an exact duplicate butwithout the lines that we are interested in, as shown in the top panel of Figure 3.9. Thisallows us to detect lines that are contaminated by lines of other elements (superposedor blended) at our resolution. Lines don’t need to be fully isolated, as it’s possible todeblend lines when we measure the EW s, but we just need to check that a weak line isnot blended and effectively lost in the wing of a strong and saturated line. By makingan identical synthetic spectrum but without the lines we are interested in, we can checkwhat the spectrum look like without those lines. If it goes to zero at the position of thoselines, then there is nothing known that directly contaminates the line. If there is a largeresidual, then we have to reject the line. Of course there is always the possibility thatthere is something unknown contaminating our lines, and having many lines to deter-mine the abundance of an element can protect us against this unknown. By looking atArcturus spectrum, (bottom panel of Figure 3.9) we can judge if our synthetic spectraare realistic and that there are not many unknown features.

Now that we have the means to determine accurate abundances, we will, in chapter 4,explain in detail how to apply this theory into practice with our FLAMES data set.


Figure 3.9: (top) Synthetic spectra of a star with stellar parameters similar to Arcturus(Teff = 4250, [Fe/H] = 0.5, log g = 1.5). The thin line contains all known atomic andmolecular lines. The thick one has the same but without the lines present in our linelist. The resolution of this spectrum is R ≈ 20 000. (bottom) UVES archive spectrumof Arcturus at a resolution of R = 120 000. The Fe i line in the centre and Si i line onthe left were kept for the analysis, despite the Fe i line not being fully isolated at ourresolution. The Fe ii line right next to it was rejected since it is a residual flux directlyunderneath it. The strong V i on the right was also considered usable, but flagged as“unsure”, due to the proximity of another potentially strong line.

Chapter 4Abundances with the FLAMESmulti-fibre instrument

As described in chapter 3, the spectroscopic determination of chemical abundancesin individual stars using atomic absorption lines can be obtained by using stellar

atmosphere models to predict the abundance of an element based on the relative depthsof absorption lines. The most accurate way to determine the iron abundance of a star isto directly measure the strength of as many iron lines as possible. This is typically doneon a spectrum of high resolution (HR), with a resolving power R ∼ 40 000 (∼ 0.125 Å atλ = 5000 Å). For precision, it is necessary to have as many lines as possible, therefore alarge wavelength coverage (∼ 2000 Å) is required. These are the optimal conditions fora detailed abundance analysis of individual stars. In this chapter, I will summarise theadded complexity of making such an analysis with an instrument (FLAMES) that hasonly half the resolution (R ∼ 20 000), a smaller wavelength coverage but a high degreeof multiplexing that requires the automation of many tasks usually carried out star bystar. Our FLAMES observations include the three Fornax stars observed by Shetroneet al. (2003) with UVES. In this chapter and in chapter 6, we will use them as a referencepoint for comparison and will refer to them as the (three) Shetrone stars. Presented inTable 4.1 are the corresponding ID from our work & their work.

Table 4.1: Name and coordinates of the 3 Shetrone stars.Our ID Shetrone RA(J2000) DEC(J2000)BL239 Fnx-M25 02 39 47.09 -34 31 49.8BL266 Fnx-M12 02 40 10.00 -34 29 58.8BL278 Fnx-M21 02 40 04.38 -34 27 11.3

48 chapter 4: Abundances with the FLAMES multi-fibre instrument

Figure 4.1: (top): Comparison of two spectra of the same star (BL239) at differentresolutions. The spectra are normalised to have their continuum flux equal to 1, asindicated by the horizontal line. The GIRAFFE (thick line) spectrum has a lower reso-lution compared to UVES (thin line). The dashed line is placed on top of an Fe i line ofmedium strength (EW ≈ 100 mÅ.) (bottom): Same as top panel but the UVES spectrumis convolved to the GIRAFFE resolution.

4.1 UVES vs FLAMES

In this thesis, two HR instruments have been used: UVES (chapter 5) and FLAMES(chapter 6). UVES is the type of instrument that is classically used to make high res-olution abundance analysis of individual stars; it has a high resolution over a largewavelength range. UVES typically looks at one star per exposure, whereas FLAMES/GIRAFFE∗, allow us to observe 100+ stars per exposure but with a factor two lowerin resolution, as illustrated in Figure 4.1 (top panel) where we show two spectra of thesame star. The GIRAFFE spectrum is from our sample and the UVES spectrum is fromthe work of Shetrone et al. (2003), convolved to the GIRAFFE resolution (bottom panel)for visual comparison.

∗ GIRAFFE is the dedicated intermediate resolution spectrograph on FLAMES

4.1: UVES vs FLAMES 49

4.1.1 UVES

The Ultraviolet and Visual Echelle Spectrograph (UVES) has two arms, the blue (UVto blue) and the red (visual to red). For our observations of individual stars in theFornax globular clusters (chapter 5) we used the red arm of UVES, covering wavelengthsbetween 4800 − 6800 Å. Shetrone et al. (2003) used exactly the same setup for thethree comparison stars. We used a slit of 1-arcsec, giving us a resolving power of about40 000. The different orders of the echelle spectra are dispersed onto two CCDs. Moreinformation on UVES can be found in Dekker et al. (2000). Nothing more will be saidabout UVES in this chapter, (more details are in chapter 5) but it will be used as a pointof reference for the FLAMES analysis, to compare our abundance results to a “classical”reference.

4.1.2 FLAMES

The Fibre Large Array Multi Element Spectrograph (FLAMES) is a fibre-fed, multi-object instrument connected to GIRAFFE, a spectrograph which has a medium-highresolution (R = 7500− 30 000). GIRAFFE is not an acronym, the name comes from theoriginal design concept, where it was standing vertically on a platform. It can cover theentire visible range 3700− 9000 Å, but not at the same time. The higher the resolutionthe shorter the wavelength coverage. It has two gratings, one low (LR) and one high-resolution (HR). On a given exposure, only a single order can be observed, referred to asa setup. In total, there are 24 HR setups and 8 LR setups. Different observing modes areavailable, ARGUS, Integral Field Unit (IFU) and MEDUSA. We only used the MEDUSAmode, which consists of up to 132 individual fibres that can be “placed” on the sky. Eachindividual fibre has a footprint (aperture) of 1.2 arcsec on the sky. The resulting spectraare projected onto a CCD with a spatial scale of 0.3 arcsec/pixel. More information aboutFLAMES is available in Pasquini et al. (2002). For our observations of the Fornax fieldstars, (chapter 6) we used FLAMES/GIRAFFE in MEDUSA mode with setups HR10,HR13 and HR14. In Table 4.2 we give an overview of the different setups used, with theold/new definition for HR14 referring for observations taken before/after October 10th

2003. On that date, HR14 grating was changed for a higher-efficiency version.

Table 4.2: Wavelength coverage and resolution of FLAMES/GIRAFFE in MEDUSAmode for setups HR10, HR13 and HR14. (Effective coverage may be slightly smaller).

setup λmin(Å) λmax(Å) RHR10 5339 5619 19800HR13 6120 6406 22500HR14 (old) 6383 6626 28800HR14 (new) 6308 6701 17740


4.2 The FLAMES Spectra

4.2.1 Extracting, calibratingWe used a pipeline made by the Geneva Observatory, girBLDRS∗ (GIRAFFE Base-LineData Reduction Software) to extract and calibrate our spectra, except for the subtractionof the sky emission lines and continuum, for which we used a software written by MikeIrwin. A master sky spectrum is made out of the 10-20 sky spectra taken during the ob-servations, is split into line and continuum components that are scaled and aligned beforebeing subtracted to the object spectrum (more details in Battaglia et al. in prep). Thespectral information needs to be extracted taking into account that the light does notnecessarily follow a straight line on a CCD row or column and the pixel/wavelength ratiovaries with wavelength. There is more than one spectrum on the CCD, correspondingeither to another star, an “empty sky” or a calibration lamp. There are five fibres tar-getted on a calibration lamp (Thorium-Argon) for simultaneous wavelength calibration,which are used for the zero-point correction with respect to the daytime wavelength andto adjust the transverse PSF used in the “optimal” extraction method. The flat fieldingis done on the extracted spectra (NFF, narrow flat fielding). After the extraction, thespectra can be rebinned into a constant wavelength increment per pixel.

4.2.2 CombiningOur data consist of multiple exposures of the same stars in three setups, as can be seenin Table 4.3. To reach the signal to noise required for our analysis, all the spectra of eachstar need to be stacked together. Before co-adding, the spectra need to be on the samerest frame and since our observations were made during two distinct periods, Septem-ber 2003 and January 2004, the observed radial velocities (Vrad) were different. Ourspectra were corrected for the heliocentric motion before they were co-added. Detailsof how we determined the Vrad and the heliocentric correction of our stars are availablein section 4.2.3. Within each observation run, the heliocentric corrections varied verylittle, by amounts typically smaller or of the same order as the uncertainty on the derivedradial velocities (∼0.5 km/s), and inducing a completely negligible shift compared to thetypical line width in GIRAFFE (FWHM ∼15 km/s).

The combining of the heliocentric corrected spectra was carried out with the IRAF†

task scombine, using a flux weighted average with median sigma clipping (cosmic rayremoval) for each period, in order to create two master spectra, one for September andone for January. These two stacked spectra were then combined in a flux weightedaverage. For the two HR14 set of spectra with different resolution and coverage, we usedonly the common overlapping section, (∼6400-6600 A) convolved the higher resolutionHR14 (old) to the resolution of HR14 (new) and treated them like our other spectra.Figure 4.2 show the extracted spectra of the three HR setups used in our observations,HR10, HR13 and HR14. By looking at the top panel of Figure 3.5, we can verify thatthe individual Fe i abundance of this star are scattered around the same central value,meaning that the three continuum levels are compatible.

∗ http://girbldrs.sourceforge.net/† http://iraf.noao.edu/

http://girbldrs.sourceforge.net/

http://iraf.noao.edu/

4.2: The FLAMES Spectra 51

Figure 4.2: Final spectra of star BL239, in each of the three GIRAFFE setups, HR10,HR13 and HR14. Overplotted on it is the continuum that DAOSPEC used for the EWmeasurement.


Table 4.3: FLAMES Exposure time logdate Exposure time (s)

HR10 HR13 HR142003-09-29 0 14400 62252003-09-30 3600 0 141022003-10-01 10800 0 69002004-01-14 3600 0 02004-01-15 0 3600 02004-01-19 0 7200 02004-01-20 0 3600 02004-01-21 0 0 72002004-01-22 0 0 75032004-01-23 3600 0 36002004-01-24 3600 0 02004-01-26 3600 0 0Total 8h 8h 12h39m

4.2.3 Determining the radial velocities (Vrad)In total, we have six independent measurements of the Vrad per star, one per setup (HR10,HR13 and HR14) and one per period (September and January). Listed in Table 4.4 arethe final Vrad(weighted average of the six values) with their associated error (the standarddeviation of the six). These measurements were made with the girBLDRS routine calledgiCrossC.py, which takes a template spectrum (G2-type star in our case) and cross-correlates it with each observed star. There were no systematic difference from setup tosetup, nor for epoch to epoch. Our mean heliocentric velocity (Vrad) is 55.9 km/s witha line of sight velocity dispersion σ = 14.2. The typical (median) error we have on avelocity (see Table 4.4) is ' 0.55 km/s. We can compare our values to the larger sampleof Battaglia et al. (2006), with a mean heliocentric velocity Vrad = 54.1 ± 0.5 and a lineof sight velocity dispersion of σ = 11.4 ± 0.4. Figure 4.3 shows the histogram of thedistribution of Vrad, with the central velocity and the 3σ membership cut-off, in solidlines (this work) in dashed lines for Battaglia et al. (2006) values. Only one star is aclear non member (BL109), with a negative Vrad, and three others are just on the edge ofmembership. These three stars have been analysed as if they were confirmed members,since they were on the edge of our membership cut-off. Our subsequent analysis showsthem to be consistent with RGB stars and not foreground dwarf stars.

4.2.4 Measuring the Equivalent WidthsEW s are classically and arguably most reliably measured line by line by hand usingSPLOT in IRAF. This can also be automated with DAOSPEC∗, a software tool that wasoptimised to work with GIRAFFE HR spectra. DAOSPEC fits a continuum over an entirespectrum, subtracts it and then measures the EW s for all the lines. It will also findthe radial velocity of the star by cross-correlating the spectrum with a line list. In both∗ http://cadcwww.dao.nrc.ca/stetson/daospec/

http://cadcwww.dao.nrc.ca/stetson/daospec/


Table 4.4: Vrad and associated errors for our Fornax targets. BOLD are possibleforeground stars, since they fall outside the accepted range of Vrad for membership inFornax. Struck-out stars have been rejected because their September average velocity wassignificantly different (>3 km/s) to their January velocity. This suggests the possibilitythat they are binary stars and they were therefore rejected from our analysis at thispoint.

Star Vrad Stddev Star Vrad Stddev Star Vrad StddevBL038 62.86 0.74 BL147 52.39 0.61 BL221 54.25 0.32BL045 79.54 0.23 BL148 65.48 0.65 BL227 52.15 0.67BL052 30.32 1.25 BL149 50.70 0.40 BL228 55.00 0.34BL061 82.38 1.12 BL150 68.70 1.12 BL229 63.09 0.57BL065 63.86 0.78 BL151 41.41 0.58 BL231 63.28 1.11BL069 68.85 2.63 BL153 50.68 0.53 BL233 58.91 0.40BL070 60.43 2.85 BL155 46.18 0.21 BL239 47.47 0.35BL076 52.83 0.78 BL156 30.63 0.55 BL242 39.48 0.55BL077 56.23 1.49 BL158 41.26 0.67 BL247 40.40 0.63BL079 56.31 0.56 BL160 53.50 0.70 BL249 63.30 0.73BL081 50.58 0.22 BL163 41.29 0.48 BL250 42.06 0.36BL084 58.92 0.80 BL165 70.01 0.71 BL251 55.14 0.45BL085 88.97 1.17 BL166 57.36 0.51 BL253 67.48 0.39BL091 52.79 0.44 BL168 75.50 0.78 BL254 73.05 0.26BL092 53.99 0.37 BL171 58.27 1.19 BL257 56.51 0.36BL094 47.45 0.43 BL173 44.93 1.60 BL258 49.41 0.79BL096 72.90 0.37 BL180 72.51 0.21 BL260 61.11 0.70BL097 47.47 0.18 BL181 57.46 1.89 BL261 71.35 1.27BL100 55.92 0.46 BL183 19.09 0.32 BL262 47.72 0.52BL104 51.36 0.31 BL185 55.32 0.44 BL266 52.47 1.43BL107 63.97 0.42 BL189 82.78 0.47 BL267 59.04 0.48BL109 -15.20 0.42 BL190 57.31 0.24 BL269 45.33 0.30BL112 53.90 2.36 BL195 39.03 0.46 BL273 63.77 6.18BL113 59.21 0.31 BL196 79.09 0.25 BL274 41.22 1.65BL115 70.05 0.30 BL197 46.27 1.02 BL278 52.43 0.56BL119 56.25 1.95 BL198 45.06 11.29 BL279 67.66 0.73BL122 65.88 0.75 BL203 56.40 0.92 BL293 56.83 0.99BL123 73.85 0.71 BL204 59.30 0.81 BL295 42.56 0.37BL125 58.23 0.77 BL205 59.91 0.27 BL298 62.09 1.43BL127 46.60 0.67 BL207 48.82 0.42 BL300 69.67 0.41BL132 40.34 0.48 BL208 54.35 0.53 BL304 71.21 0.26BL135 45.64 0.57 BL210 51.51 0.32 BL311 50.86 0.23BL138 56.52 0.33 BL211 50.35 0.20 BL315 53.15 0.85BL140 46.13 0.69 BL213 71.82 0.45 BL323 52.72 1.15BL141 76.47 0.35 BL216 66.92 0.26 BL325 18.70 0.80BL146 50.01 0.61 BL218 63.33 0.46


Figure 4.3:Histogram of For-nax radial velocitymeasurements, Vrad.The central value± 3σ cut-off areshown in solid lines(this work) andthe dashed lines(Battaglia et al.2006).

cases, the placement of the continuum is critical because it influences the derived EW s.In the case of the three GIRAFFE setups, a badly determined local continuum level willcreate an offset between the average value for lines of the same element measured indifferent setups. Since DAOSPEC is run independently on each of the three setups, thecontinuum levels in each setup is independently determined, so continuum placementproblems will show up as a systematic differences in the abundances deduced in differentsetups. Figure 4.2 shows where the continuum was placed for one star of our sample.

Figure 4.4, shows a comparison between the equivalent widths measured by SPLOTand DAOSPEC, on a UVES spectrum and on a GIRAFFE spectrum for the same star.The stars selected for the comparison are two of the three stars we have in commonwith Shetrone et al. (2003). There does not appear to be a correlation of the quality ofthe match between SPLOT and DAOSPEC with the signal to noise ratio (SNR), shown inthe bottom corner of each panel. The SNR have been calculated in the same region ofthe UVES CCD Lower and Upper (l,u) and on GIRAFFE HR 10, 13 and 14. There isgood agreement between the two methods, the EW s from both the UVES and GIRAFFEspectra seems to be scattered around the slope = 1 line. It can be seen that especially forEW . 200 mÅ, there is a good agreement between the two methods, which means thatunder these conditions, DAOSPEC can be used instead of SPLOT for EW measurements.


Figure 4.4: Comparison of the equivalent width measurements (mÅ) for two referencestars, BL239 (top) and BL278 (bottom) in different cases. The Y-axis is the SPLOTmeasurements while the X-axis is DAOSPEC. The left column EW measurements are fromUVES spectra while the right column EW s are from GIRAFFE. We plotted a full line(with slope = 1) for the perfect correlation, a dashed line with an offset of ±6 mÅ fromthis line and dotted lines as 10% error convolved with the 6 mÅ error. Also printed on thefigure are the mean difference between two measurements and the standard deviation.


4.2.5 Cleaning up the spectra

Telluric absorption lines

Telluric lines are absorption lines caused by molecules (O2, H2O) in our atmosphere. Wecannot predict their strength very well (since our atmosphere is not static) but we knowat which wavelength they occur (same rest frame for every exposure). Since no specialcalibration (e.g. fast rotating hot star) was acquired together with our observations, wecould not correct for telluric absorption lines. The only option we have is to flag (ouratomic) lines affected by telluric lines and remove them from our abundance analysis.This has to be done star by star, because the position of stellar lines changes with Vrad.This means that lines that are rejected for one star might be used in another star, becauseof their different radial velocities. Because our GIRAFFE observations were made at twodifferent times, (September 2003 and January 2004), twice as many lines per star can beaffected by tellurics and therefore rejected.

Leaks from Calibration Lamps

There was a problem with the calibration lamps in setup HR14 during our January 2004run. This was not a problem in our September observations, but after this the gratingof FLAMES was upgraded to a higher efficiency. Saturated light from the calibrationlamps leaked on our stellar spectra, as seen in Figure 4.5. This extra light artificiallyincreased the flux of the adjacent stellar spectra. Only a small fraction of our samplewere affected, and for these stars, all lines in the affected regions were simply removedfrom our abundance analysis.

Figure 4.5: Portion of a raw CCD image (HR14, January 2004) displaying the verticalspectra aligned side by side. The two brightest ones are the saturated calibration lamps(thorium lines), leaking onto their neighbouring spectra. This portion of spectrum isabout 50Å long, starting close to the Hα absorption line at the bottom of the figure.

4.3: Selecting our stellar parameters 57

4.3 Selecting our stellar parametersOnce we have accurately measured the radial velocities of our stars and the EW s for allour absorption lines, the data reduction is finished and we can proceed with the analysis.In this section, we will summarise how we assign stellar parameters to our results.

4.3.1 Photometric gravityWhen setting stellar parameters to an observed star, Fe ii lines can be used to constrainthe gravity of the star, by ensuring that there is ionisation equilibrium between Fe ii andFe i. Unfortunately, our Fe ii lines are too uncertain (not enough lines are available andthey have a large scatter) to be used for an accurate ionisation balance determination.Thus we used equation 3.9 to determine a photometric estimate and use this gravity asa final value, without adjusting them to balance Fe i and Fe ii. The distance modulus weused in this equation is 20.65, the mass of our stars was set to 1.2 M� and the extinctionE(B−V ) = 0.03 (Bersier 2000). A standard reddening law (A(V )/E(B−V ) = 3.24) wasadopted, and the bolometric corrections were computed for each star using the calibrationof Alonso et al. (1999b). The gravity does not significantly affect our abundance: loweringlog g by 0.5 dex will lower the Fe i by ≈ 0.1 dex and Fe ii by ≈ 0.3 dex. We thought itwas preferable to use the same log g scale for all stars rather than modifying our sampleusing a sometimes poorly determined [Fe ii/H].

4.3.2 Photometric Teff

As explained in section 3.2.1, we are using photometric colours to estimate Teff , usingthe calibrations found in Ramírez & Meléndez (2005). Using optical (V , I) data fromour photometric survey and infrared (J , H, K) data that were kindly provided to usahead of publication for 60% of our stars (Gullieuszik et al., in prep.), we calculated fourTeff based on four colours, V − I, V − J , V −H and V −K. The equation to calculatethe temperature coefficient θ′eff is the following (taken directly from Ramírez & Meléndez2005):

θ′eff = a0 + a1X + a2X2 + a3X[Fe/H] + a4[Fe/H] + a5[Fe/H]2 (4.1)

where X represents the photometric colours, V − I, V − J , V −H, V −K, ai are thecoefficients of the fit. The coefficients of the equations are given in Table 4.A1. Thisrelates to Teff in the following way:

Teff =5040θ′eff

+ P (X, [Fe/H]) (4.2)

the correction polynomial P , which is colour and metallicity dependent, is given as

P = P0 + P1X + P2X2 + P3X

3 (4.3)

where the coefficients (P0, P1, P2, P3) are given in Table 4.A2. Our three comparisonstars, for which we have UVES spectra, were used to make sure that our temperatureis on the correct scale (abundance is independent of χex). For these three stars, thespectroscopic Teff were in perfect agreement with the Teff(V − I). There was small offsetfor the temperature derived with the IR colours (Teff(V − {J,H,K})). We shifted them


by a constant to make them agree with the spectroscopic Teff . Doing this is the equiva-lent of zero-pointing our photometry, to be in agreement with the spectroscopy. This isrepresented graphically in Figure 4.6. The gray circles are for those stars for which wehave V , I, J , H and K photometry. The black diamonds are those for which we onlyhave V and I, where in this case we estimated their corresponding IR colours based onthe linear regression with their V − I colour, as indicated by the solid line. From theright hand side of Figure 4.6, we see that most estimates of Teff for the same star are thesame with a precision of ± 50 K.

The four Teff(V −X) as a function of (V −X) are shown in Figures 4.7, 4.8, 4.9, and4.10. In these figures are also plotted (large symbols) the model predicted (V − X) ofstellar atmosphere models at different temperature with different metallicity and gravity.These four plots clearly illustrate the relation of colour versus temperature. The predictedcolours from the models, at the metallicity and gravity range of our sample sit well onthe observed points, a nice self-consistency check. The models that we use to deduce ourabundances also produce the right colours for our stars. The final Teff we used is theaverage of the four Teff , and they are presented in Table 4.A3.


Figure 4.6: Teff(V − I) as a function of Teff(V −{J,H,K}) and the difference betweenthe two. The Teff(V − {J,H,K}) have been shifted onto the Teff(V − I) scale.


Figure 4.7: Teff(V − I) asa function of colour (V −I).Also plotted are the pre-dicted colours from modelsof different metallicity andgravity.

Figure 4.8: Teff(V − J) asa function of colour (V −J).Also plotted are the pre-dicted colours from modelsof different metallicity andgravity.


Figure 4.9: Teff(V −H) asa function of colour (V −H).Also plotted are the pre-dicted colours from modelsof different metallicity andgravity.

Figure 4.10: Teff(V − K)as a function of colour(V − K). Also plotted arethe predicted colours frommodels of different metallic-ity and gravity.


4.3.3 Iterating on the parametersStellar parameters were chosen for our stars in an iterative process. We first took thephotometric parameters, log g and Teff , as explained in sub-sections 4.3.1 and 4.3.2. Asa first guess, we set [Fe i/H] and vt to be -1.0 dex and 2.1 km/s (typical values for oursample) for all stars. We modified these starting values until we obtained good fit foreach parameter of the model. The vt will be correct when the slope measured in [Fe i/H]vs EW (slopeW ) becomes zero, as shown in the bottom panel of Figure 3.5. A fast way toconverge to the solution, is to take advantage of the linearity (and symmetry) of slopeW

on vt, (as will be demonstrated later in section 4.3.4, Figure 4.14) The simple correctionbased on the measured slope is the following:

vt = vt + slopeW /0.0055 (4.4)

At the same time as we converge on vt, we also adjust the metallicity of the model usedto reflect the average value of all the Fe i lines of the star. We did that for all the starsand our results are summarised in Figure 4.11. At this point, we still have our photo-metric first guesses for Teff and log g. This figure shows the quality of our model fittingto all the stars, where the top panel displays the values of the slope in the [Fe/H] vsχex plane for each star, and in the [Fe/H] vs EW plane in the bottom panel. In eachpanel, there are two plots, where the the X-axis is always the absolute value of the slopes,the parameter we want to minimise. The top part of each panel is a simple histogram,while the bottom one is a scatter plot of the (absolute) value of the slope against itserror. This is used to compare the quality of the fit from star to star, easily identifyingoutliers. These slopes (definition and relevance) are first introduced in chapter 3 (Fig-ure 3.5). The two shaded areas represent regions where: slope > σ and slope > 2σ .We have drawn two (solid) lines for the (empirically determined) maximum acceptableslope and error in order to still consider a good model fit. The stars outside these limits,those with large slopes and/or large scatter (σ) were not used in our analysis (chapter 6).

Once we were satisfied with vt and [Fe/H], we iterated on Teff , trying to minimise theslopeχ. Unlike vt, our Teff estimates are based on reliable photometry. We have seen fromFigure 4.6 that our precision on Teff is of the order of ± 50 K. This limits the amount bywhich we want to depart from the photometric value. Figure 4.12 show the difference inslopeχ when we change the temperature by 100 K (twice our photometric precision) forthe stars with large slopeχ. This allowed us to get some slopeχ closer to zero, thereforehaving a stellar model that is a better representation of the observed star. We stoppediterating at this point and accepted these parameters for our models, rejecting stars witha slopeχ > 0.06 and associated σ(slopeχ) > 0.05. For the vt, we rejected stars with aslopeW > 0.00035 and associated σ(slopeW ) > 0.00090.


Figure 4.11: Status of the slopes in χex (top) and in EW (bottom) after many iterationson vt and metallicity. The shaded areas correspond to slope > {1 and 2} σ. This is agraphical way of summarising our attempt to minimise the two slopes. The solid lineson the scatter plots are the maximum tolerated values in the slope and error.


Figure 4.12: Status of the slopes in χex (top) and in EW (bottom) after we allowed theTeff to change from the photometric value by a max of 100 K (our photometric precision).


4.3.4 Precision and error estimatesDAOSPEC gives an error estimate for each EW it measures, δEW which is used to calculatethe difference in abundance between EW and EW ± δEW . The largest error of the twois then adopted as our DAOSPEC abundance error, δDAO. The final adopted measurementerror for each element was calculated in the following way. In order to avoid really smallerror estimates due to low number statistics, we use the dispersion of Fe i as a lower limitof the dispersion for an element. For the final error on each of our [X/H] ratio, we haveadopted the maximum of these three values:

δ([X/H]) = MAX

(δDAO,

σ(Fe I)√(NX)

,σ(X)√(NX)

)(4.5)

and for the [X/Fe] ratios, we just take the quadratic sum of [X/H] and [Fe/H]. This con-servative estimate includes all sources of errors due to the measurement. There is alsothe uncertainty caused by our choice of stellar parameters, which is linked to the method,not the measurements and for this reason, is not included in our error bars but presentedseparately in Table 4.5. The dependencies on model atmosphere parameters is differentfor each element and each ionisation state. In the case of Fe, the only element used toconstrain our stellar parameters, it is also used for weights when trying to determine theslopes (Figure 3.5).

To determine the precision of our abundance results, we took one star (BL239) andmodified the model parameters many steps away from our adopted fit to check how thisaffected the derived Fe abundances. This illustrates the precision of the stellar param-eters, and by how much they can change and still be consistent. In Figure 4.13, weplotted the slopeχ as a function of the change in Teff (by steps of ± 100 K), along withthe corresponding change in [Fe i/H] (label above the points). It is clear from this figurethat the change (in abundance) is not symmetrical, as cooling the star by 100 K, 200 Kand 300 K does not significantly change [Fe i/H] while warming it up by steps of 100 Kincreases the abundance by at least 0.05 dex at each step. The values below the pointscorresponds to the the slope/σ, which is an indication of how far we are from a goodparameter fit. The linear regression and the estimation of its error was made using theroutine described in Fasano & Vio (1988).

From looking at Figure 4.13 and at the slopeχ numeric values, it could be argued thatTeff + 100 K would be a better choice than our adopted value (4123 K). Indeed, warmingup all of our Teff by 100 K could make all our slopes closer to zero but the resultant Teff

would not be in agreement with our photometry and the spectroscopically determinedTeff from the 3 Shetrone stars. This effect is probably caused by the combination of ourresolution and/or choice of lines and it means that the precision of our spectroscopic Teff

is not as good as our photometric one. An other way to check our Teff scale is to look atthe two curves of growth (CoG) of Figure 4.15. The top plot is for the model at Teff =4123 K and the bottom plot for a model at 4523 K. Although not immediately obvious,the hotter model has the tendency to separate the lines of different χex, with squares (�,low χex) being artificially higher (perpendicular to the curve of growth) than the pluses(+, high χex), compared to the cold CoG. A representative curve of growth should notshow such separation between different χex.


Figure 4.13: Changeof Teff for star BL239,where ∆Teff = 0 corre-sponds to Teff = 4123 Kand each step is 100 K.On the Y-axis is theslope in abundance ofFe i vs χex, where slope= 0 is a good modelfit for this parameter.On top of each point isthe resultant Fe i abun-dance (differential) de-rived with these mod-ified stellar parameters,and on the bottom is theslope/σ value, which isan indication of how farwe are from a good pa-rameter fit.

Figure 4.14: Change ofvt for star BL239, where∆vt = 0 corresponds tovt = 2.1 km/s, and eachstep is 0.1 km/s. Onthe Y-axis is the slopein abundance of Fe i vsEW , where slope = 0 isa good model fit for thisparameter. On top ofeach point is the resul-tant Fe i abundance (dif-ferential) derived withthese modified stellar pa-rameters, and on thebottom is the slope/σvalue, which is an indi-cation of how far we arefrom a good parameterfit.


Figure 4.15: Two curves of growth (CoG) for the same star, BL239. The upper exam-ple has the optimum stellar parameters while the lower example had its Teff increasedby +400 K, showing the different placement of the lines according to their χex on thecurve. On the hotter CoG, lines with squares (�, low χex) are artificially slightly higher(perpendicular to the curve of growth) than the pluses (+, high χex), compared to thecold CoG. This is an indication that the hotter CoG is not a perfect match for the data.


Table 4.5: Errors due to stellar parameters uncertainties.Element ∆ Teff = +200 K ∆ log g = -0.5 dex ∆ vt = +0.2 km/s Combined[Ba ii/H] -0.03 0.15 0.26 0.30[Ca i/H] -0.23 0.04 0.07 0.24[Cr i/H] -0.32 0.11 0.09 0.35[Eu ii/H] 0.03 0.22 0.04 0.22[Fe i/H] -0.04 0.12 0.08 0.15[Fe ii/H] 0.32 0.28 0.04 0.43[La ii/H] -0.05 0.23 0.04 0.23[Mg i/H] -0.06 0.00 0.06 0.08[Na i/H] -0.18 0.01 0.01 0.18[Nd ii/H] 0.00 0.20 0.02 0.20[Ni i/H] -0.03 0.14 0.05 0.15[Si i/H] 0.14 0.12 0.02 0.19[Ti i/H] -0.33 0.08 0.04 0.34[Ti ii/H] 0.10 0.21 0.04 0.24[Y ii/H] 0.06 0.19 0.02 0.20

We made the same plot for vt, illustrated in Figure 4.14. Unlike Teff , the changes invt are symmetrical, making it easy to find the right vt to represent our star. ∆vt = 0corresponds to vt = 2.1 km/s and from this plot, looking only at the value of the slopes,vt = 2.2 would also be acceptable. Both of them have slopeW � σ(slopeW ), showing thelimit of our precision, which should be of the order of 0.2 km/s.

4.4 Systematics and correctionsIn this section we discuss the systematic effects due to the method of analysis and theeffect of hyperfine structure on the Eu line at 6645.1 .

4.4.1 SystematicsTo measure systematics caused by different technical choices, we have used five differentmethods and compared them over a range of elements. We are interested in measuringthe effect of changing the stellar models, the line list, the resolution and the way ofmeasuring the EW s. We analysed the 3 Shetrone stars in five different ways, alwaysre-determining the stellar parameters to obtain a good fit:

• The original 2003 abundance analysis, using UVES spectra, the Shetrone line list,splot to measure the EW s and extrapolated (in log g and Teff), plane-parallelmodels (Gustafsson et al. 1975).

• Same as above but using spherical MARCS 2005 models instead of the plane-parallelmodels of Gustafsson et al. (1975).

4.4: Systematics and corrections 69

Figure 4.16: Abundance differences between the 3 Shetrone stars, using different meth-ods of analysis on the UVES spectra, (open symbols) compared to the GIRAFFE analysis(filled dots). We plotted the original 2003 results (inverted triangles) our re-analysis withthe 2005 models, keeping the old line list + splot (triangles), new line list + splot (dia-monds) and new line list + daospec (circle).


Figure 4.17: Curve ofgrowth for BL239 (Fnx-M25), showing a lack ofweak lines, leading to aninaccurate abundance.

• Same as above but using our new line list, better suited for the more metal rich starsin Fornax, restricting the analysis to the same wavelength coverage as FLAMESHR 10, 13 and 14.

• Same as above but measuring the EW s automatically with DAOSPEC instead ofmanually with splot.

• Same as above but instead of using the same UVES spectra like the four previouscases, we used our new, independent GIRAFFE spectra, effectively reducing theresolution by a factor two (from R ' 40 000 to R ' 20 000).

We compare the abundance determined in each these cases for 14 different elementsor ionisation states in Figure 4.16. The most striking feature in this plot is the hugedifference between different analysis (0.4 dex) in [Fe/H] for star BL239, the one with the[Fe/H] varying from -0.9 to -1.3. The reason why the GIRAFFE analysis differs so muchfrom the original UVES analysis (Shetrone et al. 2003) is due to the fact that they lackweak lines, as can be seen in Figure 4.17, where the lack of lines on the left part of theCoG is evident and can lead to bias stellar parameters and therefore abundances. In thisparticular case, the lack of weak lines drove the microturbulence to lower values, biassingthe metallicity towards lower values. The other two stars do not show a similar problembecause the most metal-poor star had enough weak lines, and Shetrone et al. themselvesadded extra weak lines to the analysis of the metal-rich one.

Old models versus new models

There are several differences between the old (Gustafsson et al. 1975) and new (Gustafs-son et al. 2003) family of models. The two most important are the different geometryused (plane-parallel versus spherical) and the physics involved determining the opacities.Abundances calculated with these two methods are similar, most of the time within theerror bars. Heiter & Eriksson (2006) studied extensively the effect of using differentgeometry in determining the abundance. We cannot directly compare to their result be-cause as we summarise in our section 3.3.2 they compare both p2005_p and s2005_p tothe fully consistent s2005_s while we are trying to compare plane-parallel and spherical

4.4: Systematics and corrections 71

models with different physics, both of them treated in the plane-parallel approximationfor the line formation code (p1975_p to s2005_p). According to their conclusion, if the“ideal” case is the fully consistent case s_s, using mixed geometry (s_p) is preferable(closer to s_s) than the full p_p case. Therefore, the effect of geometry alone is mea-surable and is partially responsible for the difference in abundance between the trianglesand inverted triangles symbols of Figure 4.16.

Choice of Lines to include in the analysis

To compare the effects of different line lists (see section 3.4), compare the triangles andthe diamonds in Figure 4.16. This is usually the biggest difference between two methods,showing the importance of having a common line list for comparing abundance results.Generally, there is a difference of 0.3 dex for [Fe i/H] in BL239 and a smaller difference(0.15 dex) for the other two stars. But as long as there are enough lines in a given linelist, the derived abundance is usually reliable. Only when the number of lines goes downthere can be a problem with the abundance, where an outlier can really affect the derivedvalue. But this has to be checked star by star, as it can be metallicity and/or signal tonoise dependent.

DAOSPEC vs SPLOT for measuring EW s

The method used to measure the EW s also affects the abundance results but only by asmall amount, in most cases well within the error estimate. This can be seen by lookingat the circles and the diamonds in the Figure 4.16. This check reinforces our confidenceusing the automatic DAOSPEC measurements allowing us to go from hand-measurementto a much faster automated processing.

UVES vs GIRAFFE, the effect of resolution and wavelength coverage

Comparing UVES and GIRAFFE results shows that even with a loss of a factor two inresolution, it is possible to determine accurate abundances. By comparing the emptycircles to the solid dots in Figure 4.16, it is clear that the results are identical, withinthe errors.

4.4.2 Hyperfine splitting correctionThe hyperfine structure is a small perturbation in the energy levels of an atoms due tothe interaction of the nuclear magnetic dipole with the magnetic field of the electron.The electron moving around the nucleus has a magnetic dipole moment, because it ischarged. Its interaction with the magnetic moment of the nucleus (due to its spin) leadsto hyperfine splitting of the energy level. This phenomenon is energy level dependent,therefore different for each line. We can correct for hyperfine splitting with spectral syn-thesis, as presented for the Eu abundance in chapter 5 and illustrated in Figure 5.4. Forour FLAMES data, we only corrected for the HFS in our Eu line at λ = 6645.1 Å. Hereis a summary of the method we used.

A line that has hyperfine structure consist of many small lines close to each otherinstead a line arising from a single energy level. These lines can be used to create a


Figure 4.18: HFScorrection (dex)for the Eu line at6645.1Å, tested on aplane-parallel modelof Teff = 4300, log g= 0.6, [Fe/H] = -1.5and vt = 1.7.

synthetic spectrum, where all these split lines will be superposed together. All the stel-lar parameters are known except the abundance of Eu, which we can modify until itrepresents the observed spectra. Ignoring HFS tends to overestimate the abundance socorrecting for HFS will typically lower the “direct EW” measured abundance.

Making synthetic spectra and taking into account the HFS information to find theabundance of a line is a good method but it is too time consuming for large sampleof stars or for elements with many lines. To measure the Eu abundance in the manyGIRAFFE stars, we rather designed an HFS correction to be applied to the abundancededuced from the EW treated as a single line. We found out that this correction, in ourstellar parameter and abundances range, depends dominantly on the strength (EW ) ofthe line, and very little on other parameters. The validity of this correction was testedin the following ranges:

• -2.0 < [Fe/H] < -1.0

• 1.7 < vt < 2.5

• 0.3 < log g < 1.2

• 4000 < Teff < 4800

The HFS correction becomes EW -dependent only and is applied after the the abun-dance has been derived from the non-corrected line. In Figure 4.18 we present thecorrection∗ for the Eu line at λ = 6645.1 Å.

Appendix 4.A Large tables

∗ Thanks to Kim Venn for calculating the correction in this range of parameters

4.A: Large tables 73

Table 4.A1: Coefficients of the giant star colour calibrationColour a0 a1 a2 a3 a4 a5

(V − I) 0.3575 +0.9069 -0.2025 +0.0395 -0.0551 -0.0061(V − J) 0.2943 +0.5604 -0.0677 +0.0179 -0.0532 -0.0088(V −H) 0.4354 +0.3405 -0.0263 -0.0012 -0.0049 -0.0027(V −K) 0.4405 +0.3272 -0.0252 -0.0016 -0.0053 -0.0040

Table 4.A2: Metallicity dependence of the Polynomial correction.Colour P0 P1 P2 P3 Validity(V − I) +0.42933 ... ... ... −0.5 ≤ [Fe/H] ≤ 0.5(V − J) -122.595 +76.4847 ... ... ′′

(V −H) -377.022 +334.733 -69.8093 ... ′′

(V −K) -72.6664 +36.5361 ... ... ′′

(V − I) -0.14180 ... ... ... −1.5 ≤ [Fe/H] ≤ −0.5(V − J) -10.3848 ... ... ... ′′

(V −H) +71.7949 -55.5383 +9.61821 ... ′′

(V −K) +86.0358 -65.4928 +10.8901 ... ′′

(V − I) +9.31011 ... ... ... −2.5 ≤ [Fe/H] ≤ −1.5(V − J) +4.18695 +13.8937 ... ... ′′

(V −H) -27.4190 +20.7082 ... ... ′′

(V −K) -6.96153 +14.3298 ... ... ′′

(V − I) -23.0514 ... ... ... −4.0 ≤ [Fe/H] ≤ −2.5(V − J) -67.7716 +28.9202 ... ... ′′

(V −H) -46.2946 +20.1061 ... ... ′′

(V −K) -943.925 +1497.64 -795.867 +138.965 ′′

Table 4.A3: Coordinates and photometric temperatures in different colours, shifted on theTeff(V − I) scale. Teff is the average of the four different TV −X .

Star RA (J2000) DEC (J2000) TV −I TV −J TV −H TV −K Teff log g BCV MBol

BL038 02 40 20.45 -34 24 00.1 3984 3976 3996 3964 3980 0.69 -0.93 -3.32BL045 02 40 07.52 -34 23 31.8 4114 4134 4128 4110 4122 0.85 -0.77 -3.07BL052 02 40 10.42 -34 25 17.6 4009 3999 3995 3984 3997 0.72 -0.91 -3.27BL061 02 39 28.59 -34 18 38.0 4316 4341 4335 4333 4331 0.84 -0.60 -3.30BL065 02 39 22.15 -34 19 40.3 4315 4340 4334 4332 4330 0.97 -0.59 -2.99BL069 02 39 29.56 -34 25 10.6 4143 3999 4033 3994 4042 0.72 -0.86 -3.29BL070 02 39 40.46 -34 19 38.8 3939 3928 3931 3935 3933 0.68 -1.00 -3.29BL076 02 39 31.49 -34 23 05.1 4088 4061 4063 4048 4065 0.83 -0.83 -3.05BL077 02 39 51.42 -34 21 20.9 4027 4024 4025 4028 4026 0.80 -0.88 -3.08BL079 02 39 19.60 -34 24 49.3 3986 3903 3922 3934 3936 0.76 -1.01 -3.08BL081 02 39 56.01 -34 24 10.5 4102 4045 4060 4042 4062 0.82 -0.84 -3.07BL084 02 38 42.96 -34 25 49.1 3972 3964 3966 3970 3968 0.72 -0.95 -3.23BL085 02 38 55.53 -34 25 36.3 4277 4299 4294 4293 4291 0.87 -0.63 -3.19

Continued on next page



BL091 02 39 04.31 -34 25 18.8 4155 4165 4163 4164 4162 0.86 -0.73 -3.09BL092 02 38 49.28 -34 24 04.9 3965 3956 3959 3963 3961 0.74 -0.96 -3.16BL093 02 38 39.87 -34 25 57.1 4068 4070 4070 4072 4070 0.90 -0.82 -2.87BL094 02 38 53.89 -34 25 06.7 3990 3984 3985 3989 3987 0.72 -0.92 -3.24BL096 02 39 14.33 -34 22 41.5 4012 4008 4009 4013 4010 0.75 -0.90 -3.19BL097 02 39 04.07 -34 23 52.7 4058 4059 4059 4062 4060 0.82 -0.84 -3.06BL100 02 38 56.00 -34 24 44.8 4043 4042 4043 4046 4044 0.84 -0.86 -3.01BL104 02 39 14.59 -34 23 21.0 4014 4011 4012 4015 4013 0.77 -0.89 -3.15BL107 02 38 54.37 -34 31 23.5 4366 4396 4389 4387 4384 0.83 -0.56 -3.37BL109 02 39 04.08 -34 37 58.7 4334 4360 4354 4352 4350 0.93 -0.58 -3.09BL112 02 38 45.68 -34 34 47.5 4096 4100 4099 4101 4099 0.76 -0.79 -3.27BL113 02 39 08.16 -34 36 53.3 4179 4191 4188 4189 4187 0.83 -0.72 -3.19BL115 02 38 43.45 -34 32 05.3 4112 4117 4116 4118 4116 0.79 -0.77 -3.22BL119 02 38 42.17 -34 29 50.8 3964 3956 3959 3963 3960 0.73 -0.96 -3.19BL122 02 39 02.91 -34 31 12.0 4051 4051 4052 4054 4052 0.76 -0.84 -3.22BL123 02 38 53.75 -34 30 06.6 3995 3990 3991 3995 3993 0.71 -0.92 -3.27BL125 02 39 08.50 -34 30 55.4 4078 4080 4080 4082 4080 0.79 -0.82 -3.17BL127 02 39 04.00 -34 37 26.1 3963 3955 3957 3961 3959 0.71 -0.96 -3.25BL132 02 38 53.62 -34 33 04.5 3916 3903 3906 3911 3909 0.65 -1.04 -3.33BL135 02 39 01.55 -34 36 48.7 4057 4058 4058 4061 4058 0.83 -0.84 -3.06BL138 02 39 16.20 -34 37 00.2 3989 3906 3924 3937 3939 0.71 -0.99 -3.21BL140 02 39 11.56 -34 30 44.7 3995 3989 3991 3994 3992 0.75 -0.92 -3.17BL141 02 39 10.99 -34 28 34.4 4076 4078 4078 4080 4078 0.84 -0.82 -3.04BL146 02 38 41.76 -34 28 58.4 4074 4076 4076 4078 4076 0.84 -0.82 -3.04BL147 02 38 44.02 -34 30 51.8 4186 4199 4196 4197 4194 0.94 -0.70 -2.92BL148 02 39 11.05 -34 39 08.6 3929 3917 3920 3925 3923 0.72 -1.03 -3.16BL149 02 38 57.19 -34 35 39.8 4097 4101 4100 4102 4100 0.88 -0.80 -2.97BL150 02 39 13.94 -34 28 36.1 4026 4023 4024 4027 4025 0.80 -0.88 -3.08BL151 02 39 00.29 -34 30 30.0 4025 4022 4023 4026 4024 0.81 -0.88 -3.06BL153 02 39 08.35 -34 32 44.8 4035 4033 4034 4037 4035 0.82 -0.86 -3.04BL155 02 38 40.44 -34 35 25.0 4059 4060 4060 4063 4060 0.90 -0.84 -2.87BL156 02 39 05.62 -34 26 30.6 4096 4100 4099 4101 4099 0.89 -0.79 -2.95BL157 02 39 02.44 -34 27 35.4 4073 4075 4075 4077 4075 0.88 -0.82 -2.93BL158 02 39 06.84 -34 35 45.7 4076 4079 4078 4081 4078 0.87 -0.82 -2.96BL160 02 39 01.74 -34 27 19.2 4027 4025 4026 4029 4027 0.84 -0.88 -3.00BL163 02 38 43.54 -34 32 25.4 4025 4022 4023 4026 4024 0.88 -0.88 -2.89BL165 02 39 13.37 -34 40 15.8 3967 3959 3961 3965 3963 0.77 -0.95 -3.08BL166 02 39 16.01 -34 30 12.9 4070 4091 4074 4108 4086 0.84 -0.81 -3.06BL168 02 39 14.78 -34 27 29.8 4014 4011 4012 4015 4013 0.83 -0.89 -3.00BL171 02 39 15.48 -34 30 46.5 4047 4047 4047 4050 4048 0.87 -0.85 -2.94BL173 02 39 14.36 -34 34 42.3 3991 3985 3987 3990 3988 0.85 -0.93 -2.92BL180 02 39 33.38 -34 33 58.4 4115 4102 4107 4134 4114 0.77 -0.78 -3.25BL181 02 39 16.76 -34 34 49.9 4016 3962 3969 3969 3979 0.80 -0.93 -3.04BL183 02 39 47.94 -34 26 43.4 4073 4199 4223 4230 4181 0.83 -0.72 -3.16BL185 02 39 27.67 -34 37 48.5 4026 3980 3999 3988 3998 0.69 -0.92 -3.32BL189 02 39 38.24 -34 31 20.3 4167 4222 4272 4274 4234 0.91 -0.67 -3.03BL190 02 39 17.80 -34 30 56.8 4004 3954 3967 3990 3979 0.82 -0.94 -2.99BL195 02 39 20.21 -34 31 57.0 4153 4153 4147 4191 4161 0.91 -0.73 -2.95BL196 02 39 34.07 -34 33 33.2 4011 4010 4010 4029 4015 0.76 -0.89 -3.17BL197 02 39 29.33 -34 26 35.5 3958 3953 3963 3948 3956 0.68 -0.97 -3.30




BL198 02 39 19.83 -34 26 27.1 4077 3988 4000 3992 4014 0.78 -0.89 -3.11BL203 02 39 50.64 -34 26 47.0 4043 4035 4046 4025 4037 0.79 -0.87 -3.13BL204 02 39 16.23 -34 32 29.9 4076 4184 4161 4134 4139 0.97 -0.75 -2.78BL205 02 39 38.06 -34 37 06.2 4229 4237 4263 4244 4243 0.96 -0.67 -2.92BL207 02 39 26.29 -34 29 03.4 3988 3979 3977 3992 3984 0.74 -0.93 -3.19BL208 02 39 32.94 -34 32 06.2 4140 4166 4161 4170 4159 0.89 -0.74 -3.00BL210 02 39 47.65 -34 27 05.4 4057 4064 4072 4057 4062 0.81 -0.84 -3.10BL211 02 39 51.18 -34 29 58.6 3940 3961 3964 3966 3958 0.69 -0.97 -3.28BL213 02 39 50.18 -34 35 59.3 4039 4045 4017 4025 4032 0.78 -0.87 -3.14BL216 02 39 41.92 -34 30 35.9 3970 4008 3995 4019 3998 0.75 -0.92 -3.18BL218 02 39 50.77 -34 28 36.5 3920 3943 3941 3953 3939 0.67 -1.00 -3.31BL221 02 39 31.85 -34 29 19.9 4051 4051 4057 4064 4056 0.81 -0.84 -3.08BL227 02 39 45.25 -34 31 57.8 4020 4064 4043 4055 4046 0.84 -0.85 -3.01BL228 02 39 53.84 -34 29 56.4 3975 3998 3990 4007 3992 0.71 -0.92 -3.28BL229 02 39 24.80 -34 34 38.1 4022 4008 4012 4013 4014 0.80 -0.90 -3.07BL231 02 39 18.96 -34 26 43.9 4087 4033 4040 4052 4053 0.86 -0.84 -2.96BL233 02 39 53.58 -34 37 50.1 4064 4028 4039 4063 4048 0.83 -0.86 -3.03BL239 02 39 47.09 -34 31 49.8 4083 4145 4127 4137 4123 0.89 -0.77 -2.96BL242 02 39 28.03 -34 34 01.2 4064 4058 4057 4074 4063 0.85 -0.83 -3.00BL247 02 39 43.07 -34 40 18.4 4050 4031 4018 4027 4032 0.85 -0.87 -2.98BL249 02 39 54.24 -34 35 11.2 4032 4034 4009 4019 4024 0.82 -0.88 -3.03BL250 02 39 45.03 -34 40 11.1 3866 3829 3849 3833 3844 0.65 -1.16 -3.27BL251 02 39 30.81 -34 35 45.1 3998 3970 3977 3985 3982 0.81 -0.93 -3.00BL253 02 39 34.08 -34 33 09.6 3989 3995 4006 4022 4003 0.82 -0.91 -3.02BL254 02 39 41.78 -34 34 16.4 4025 4013 3990 4017 4011 0.83 -0.89 -3.01BL257 02 39 57.30 -34 31 20.8 3981 3984 3994 4015 3994 0.78 -0.93 -3.10BL258 02 39 49.68 -34 28 50.4 4008 4045 4026 4039 4030 0.85 -0.88 -2.96BL260 02 39 55.38 -34 29 54.6 3985 4026 4016 4010 4009 0.79 -0.90 -3.09BL261 02 39 57.80 -34 26 48.8 4028 4053 4066 4039 4046 0.82 -0.86 -3.06BL262 02 39 38.42 -34 26 10.3 4054 4021 4027 4018 4030 0.84 -0.88 -2.98BL266 02 40 10.00 -34 29 58.8 4157 4241 4225 4223 4212 0.83 -0.68 -3.20BL267 02 40 17.50 -34 26 06.1 4197 4210 4211 4186 4201 0.80 -0.70 -3.27BL269 02 39 58.20 -34 32 05.3 3965 3992 4008 3997 3990 0.75 -0.92 -3.18BL273 02 40 09.37 -34 36 17.1 4123 4158 4124 4149 4138 0.81 -0.76 -3.18BL274 02 40 06.14 -34 28 52.0 4011 4041 4032 4019 4026 0.72 -0.87 -3.29BL278 02 40 04.38 -34 27 11.3 3964 3967 3977 3980 3972 0.64 -0.95 -3.43BL279 02 40 02.70 -34 38 29.9 4228 4298 4274 4286 4272 0.95 -0.64 -2.97BL293 02 40 01.77 -34 27 47.9 4030 4031 4024 4026 4028 0.72 -0.87 -3.29BL295 02 40 26.72 -34 26 56.8 3994 3976 3975 3975 3980 0.70 -0.94 -3.28BL298 02 40 13.49 -34 30 02.0 4118 4145 4129 4125 4129 0.85 -0.76 -3.07BL300 02 40 17.90 -34 27 00.7 3989 3991 4001 3977 3990 0.71 -0.92 -3.27BL304 02 40 05.49 -34 32 42.7 3959 3936 3941 3965 3950 0.70 -0.97 -3.26BL311 02 40 22.64 -34 31 31.0 4032 4022 4022 4032 4027 0.79 -0.88 -3.12BL315 02 40 24.22 -34 26 20.0 4173 4135 4115 4133 4139 0.86 -0.76 -3.06BL323 02 40 16.76 -34 29 34.4 3897 3859 3881 3887 3881 0.66 -1.08 -3.28BL325 02 40 27.00 -34 26 44.1 4082 4064 4036 4050 4058 0.81 -0.84 -3.09

Chapter 5High resolution spectroscopy inFornax Globular Clusters

Published as A&A 2006 453 547VLT/UVES spectroscopy of individual stars in three globular clusters in

the Fornax dwarf spheroidal galaxy∗

B. Letarte, V. Hill, P. Jablonka, E. Tolstoy, P. François, G. Meylan

ABSTRACT– We present a high resolution (R∼ 43 000) abundanceanalysis of a total of nine stars in three of the five globular clusters as-sociated with the nearby Fornax dwarf spheroidal galaxy. These threeclusters (1, 2 and 3) trace the oldest, most metal-poor stellar popula-tions in Fornax. We determine abundances of O, Mg, Ca, Ti, Cr, Mn,Fe, Ni, Zn, Y, Ba, Nd and Eu in most of these stars, and for somestars also Mn and La. We demonstrate that classical indirect methods(isochrone fitting and integrated spectra) of metallicity determinationlead to values of [Fe/H] which are 0.3 to 0.5 dex too high, and thatthis is primarily due to the underlying reference calibration typicallyused by these studies. We show that Cluster 1, with [Fe /H]=−2.5,now holds the record for the lowest metallicity globular cluster. Wealso measure an over-abundance of Eu in Cluster 3 stars that has onlybeen previously detected in a subgroup of stars in M15. We find thatthe Fornax globular cluster properties are a global match to what isfound in their Galactic counterparts; including deep mixing abundancepatterns in two stars. We conclude that at the epoch of formation ofglobular clusters both the Milky Way and the Fornax dwarf spheroidalgalaxy shared the same initial conditions, presumably pre-enriched bythe same processes, with identical nucleosynthesis patterns.

∗ Based on UVES observations collected at the European Southern Observatory, proposal number70.B-0775

78 chapter 5: HR spectroscopy in Fornax Globular Clusters

5.1 Introduction

It is now established that some dwarf galaxies have globular cluster systems aroundthem (Lotz et al. 2004, van den Bergh 2006, Seth et al. 2004). Their possible com-

mon origin with clusters in larger parent galaxies, the link between the dwarf galaxy fieldand globular cluster stars are open questions to be addressed. The largest samples ofdwarf galaxies with globular cluster systems are however distant, and this restricts theanalyses to using integrated properties.

Fornax and Sagittarius are the nearest dwarf spheroidal galaxies (dSph) with globularclusters and can be resolved into individual stars. The Fornax dSph contains five starclusters (Shapley 1938; Hodge 1961) and while the Sagittarius dSph is obscured by dustand confused by merging with our Galaxy, Fornax is high above the Galactic plane,therefore offering a uniquely useful target for investigation, see Figure 5.1.

Figure 5.1: A ≈ 85′ × 62′ DSS image of the Fornax dSph. North is up and East is tothe left, as indicated. We have marked the position of the 5 GCs using the numberingscheme defined by Shapley 1938 and Hodge 1961

The ages of the Fornax globular clusters have been determined by fitting isochronesto deep HST Colour-Magnitude Diagrams [CMDs] (Buonanno et al. 1998, 1999). Theyare found to be the same age as old metal-poor Galactic clusters M92 and M68 (around13 Gyr old) to within ± 1 Gyr, with the exception of Cluster 4, which seems buried in the

5.2: Observations 79

center of Fornax and maybe younger by about 3 Gyr. The cluster metallicities have beenestimated with different techniques ranging from fitting a slope to the Red Giant Branch(RGB) to high and medium resolution spectroscopy of the integrated light of the cluster.Conclusions vary from one work to another, as summarized in Strader et al. (2003), butthe clusters definitely appear more metal-poor than the bulk of the galaxy field stellarpopulation, with bluer RGBs, well populated blue horizontal branches (HB) and a rangeof HB morphology (Buonanno et al. 1998 and 1999). Saviane et al. (2000) showed thatthe Fornax dSph field star colour distribution is well fitted by two Gaussian functions,best interpreted as a bi-modal metallicity distribution, with the older population havinga wide abundance range between −2.2 and −1.4. Stars as young as 108 Myr have alsobeen discovered in the field of Fornax (Stetson et al. 1998). In this framework, theglobular clusters of Fornax dSph trace the first stages of star formation in the galaxy.

High resolution spectroscopy of individual stars in the clusters is the only way toassess the abundances of individual chemical species. Alpha, iron-peak, heavy -elementsprovide essential clues on (i) the conditions of formation of the globular clusters in a dwarfgalaxy, including epoch and time scales (ii) to probe the nucleosynthesis in a galacticsystem with a star formation history that is fundamentally different from that of theMilky Way. We present here a VLT/UVES spectroscopic analysis of a total sample ofnine stars in three Fornax dSph clusters.

5.2 ObservationsWe targeted Cluster 1, Cluster 2 and Cluster 3, to span the Fornax globular clustersystem range of distances from the galaxy centre avoiding regions of heavy crowding.We also sampled a range of HB morphology, as well as the metallicity and concentrationranges. Cluster 1, at a radial distance of 43 arcmin (or 1.75 kpc at the distance of FornaxdSph) from the galaxy center, is diffuse, with low surface brightness, most of its HB isred. Cluster 2, located at 25 arcmin (1 kpc) from the galaxy center, is slightly moreconcentrated and exhibits a more extended HB. Finally, Cluster 3 at a galactocentricradial distance of 13 arcmin (530 pc) is very dense and has an extended HB.

We used the red arm of UT2/UVES, CD#3, centered at 580nm, with a wavelengthrange of 480-680nm (Dekker et al. 2000) in visitor mode in October 2002. We obtainedspectra with a resolution of ∼43 000 and average S/N ∼ 20 − 30 per pixel with anintegration time of 2 − 6 hours for each of the nine individual stars in Fornax dSphglobular Clusters 1, 2 and 3. The stars were selected to be on the RGB from CMDs,(Buonanno et al. 1985; Demers et al. 1990; Jorgensen & Jimenez 1997 and Buonannoet al. 1998). Their individual finding charts are shown in Figure 5.2. We also observed5 calibration red giant branch stars in the well studied globular cluster M15 (Snedenet al. 1997). The observations of M15 stars provide an independent check on our datareduction and analysis methods. Details of the observations are shown in Table 5.1,including the derived radial velocities Vrad and S/N ratios.

80chapter

5:H

Rspectroscopy

inFornax

Globular

Clusters

Table 5.1: Observation LogOur id Litterature IDs RA (J2000) DEC (J2000) exp. time S/N @ Vrad Comments

degrees degrees (s) 670 nm (km/s)Cl1-D56 D56b, J24c 39.254958 -33.17790 14400 23 57.6 same slit as D68Cl1-D68 B51a, D68b, J23c 39.254609 -34.17875 14400 50 60.2 same slit as D56Cl1-D164 B18a, D164b, J65c, B713d 39.257554 -34.18628 18000 30 60.0Cl2-B71 B71a 39.677388 -34.80412 21600 30 63.4Cl2-B74 B74a 39.684917 -34.80301 14400 ... ... Too faintCl2-B77 B77a 39.685203 -34.80303 14400 30 64.1 same slit as B74Cl2-B200 B200d ... ... 3900 ... ... carbon starCl2-B226 B226d 39.682143 -34.80801 7200 40 64.0Cl3-B59 B59a, J9c 39.942803 -34.25855 10800 30 59.7Cl3-B61 B61a, J31c 39.957271 -34.25782 21600 30 63.7Cl3-B82 B82a, J3c 39.951057 -34.25277 14400 40 64.8M15-S1 S1e, K431f 322.484344 12.21002 900 113 -106.4M15-S3 S3e, K387f 322.481920 12.21231 1200 115 -111.3M15-S4 S4e, K825f 322.509599 12.18986 750 122 -101.4 spec. double starM15-S6 S6e, K1040f 322.543661 12.16832 1000 96 -100.0M15-S7 S7e, K146f 322.457798 12.13489 900 83 -100.7IDa from Buonanno et al. (1985) IDc from Jorgensen & Jimenez (1997) IDe from Sandage (1970)IDb from Demers et al. (1990) IDd from Buonanno et al. (1998) IDf from Kustner (1921)

5.3: Data Reduction and Analysis 81

Figure 5.2: The finding charts for our observations of the Fornax GCs, from 1 (left) to3 (right). North is up and East is left, as indicated. Note that star Cl3-B59 is outside ofthe cluster 3 HST field, to the west.

5.3 Data Reduction and AnalysisThe spectra were extracted with the standard UVES pipeline, except for two pairs ofstars on the same slit which we had to reduce interactively using the UVES contextwithin MIDAS (see Table 5.1). At the telescope we already identified Cl2-B200 as acarbon star, and it was discarded from further analysis. As already noted by Sneden etal. (1997), M15-S4 is probably a spectroscopic double star, as all lines are significantlywider (larger Full Width Half Maximum [FWHM]) than the other stars of M15. It wasnot used for our abundance analysis.

For each of our targets we made equivalent width (EW) measurements with SPLOTin IRAF, except for the lines with a small EW (. 50 mÅ). For these weak lines, wenoticed that SPLOT was giving very unstable FWHM measurements. A home-madegaussian-fitting program was used for these lines to fix the FWHM at the instrumentalvalue. We also used DAOSPEC∗, a new programme that automatically measures EWsby iteratively fitting gaussians of fixed FWHM to all lines in the spectrum and removingthe continuum signature (Stetson & Pancino, in preparation). Having confirmed thatDAOSPEC gives results compatible with those obtained by hand for lines of moderatestrength (EW ≤ 150 mÅ)†, we used DAOSPEC measured EWs for M15.

We can detect a range of elements in our coadded spectra: Fe i, Fe ii, Ti i, Ti ii,O i, Mg i, Ca i, Cr i, Mn i, Ni i, Zn i, Y ii, Ba ii, La ii, Nd ii and Eu ii, which allowsus to achieve a comprehensive abundance analysis. The most important is Fe, with anaverage of 50 measured lines for Fe i and 10 lines for Fe ii. Line parameters and EWmeasurements for all stars are reported in Table 5.A1. Abundances for the different∗ http://cadcwww.dao.nrc.ca/stetson/daospec/† We note here for completeness that, at this high resolution, the fixed FWHM gaussian hypothesis

adopted by DAOSPEC does not hold for the strongest lines (EW > 150 mÅ) where departures fromthe gaussianity and natural broadening play a significant role.

http://cadcwww.dao.nrc.ca/stetson/daospec/


elements were calculated with CALRAI, originally described in Spite (1967) with manyimprovements over the years. The stellar atmospheres models are those of Plez (privatecommunications, 2000 and 2002, described in Gustafsson et al. 1975, Plez et al. 1992,Gustafsson et al. 2003). Spectral synthesis was required for some elements: Eu, Zn, Mg,Na, O and Ba to account for hyperfine splitting (Eu, Ba); weak lines (Zn, O) and strong,possibility saturated lines (Na, Mg).

Initial guesses were made for the stellar effective temperature (Teff) using V −I and/orB − V colours, using the Alonso et al. (2001) calibration and a reddening of E(B − V )= 0.065. The surface gravity (log g) was estimated assuming a 0.8 M� mass for thestars, a distance modulus of (m-M)=20.85 mag and bolometric corrections from Alonsoet al. (2001). However, the quality of the photometric data we gathered for these starsturned out to be too poor to constrain firmly the star’s effective temperature (only 2stars had HST photometry in Buonnano et al. 1998, and the other photometric sourceswere ground-based, suffering from crowding and not all in a homogeneous photometricsystem). We therefore chose to base our analysis solely on spectroscopic criteria. TheTeff , log g and micro-turbulence velocity (vt) were adjusted to insure that we had theionisation balance of Fe i and Fe ii and that the Fe i abundance is independent of bothline strength and excitation potential of the line. Figure 5.3 illustrates the quality of our

Figure 5.3: Observed Curve of growth for Fe i in Cl3-B82. The dotted line marks the[Fe/H]=0 location, while the full line is the theoretical curve of growth for a typical Fe iline with the stellar parameters adopted for this star.

5.3: Data Reduction and Analysis 83

solution by showing the curve of growth obtained for Cl3-B82, where we notice that everypart of the curve of growth is well populated. The final set of stellar parameters used foreach star are shown in Table 5.2. The [Fe/H] in this table is the metallicity of the modelused to compute the abundances, not the final abundance value of Fe i or Fe ii of the star.

As an additional test, since the S/N reached in the individual spectra was ratherlimited, we also co-added the spectra of stars with similar parameters within each cluster(all three stars of Cluster 3 on the one hand, and the two cooler stars of Cluster 1 on theother hand), and repeated the analysis. The results are fully consistent with the analysisof the individual stars: Teff , log g and vt are undistinguishable, while the mean [Fe/H] isrecovered within 0.02 dex, and most of the other abundance ratios fall well within thestar to star scatter.

A significant source of error in our analysis is the uncertainty in measuring the EW.Our errors in the EW determinations were estimated by propagating the EW error es-timates (from splot) through the abundance computation (the abundances EW + δEWand EW − δEW were computed and compared to the central adopted value). For ele-ments which were computed by spectral synthesis, the error is estimated by eye, plottinga range of acceptable fits, as illustrated in Figure 5.4 for the weak Eu line in Cl3-B59.Another way to estimate the measurement errors affecting the abundance is to consider

Figure 5.4: The Synthetic spectra for the Eu line at λ = 6645.1 overlaid on the datafor Cl3-B59. The middle line is the adopted fit, while the lower and upper ones are theerror estimate of ± 0.1 dex. The larger line on the left is a Ni line.


the dispersion (rms) around the mean. For species with sufficient numbers of lines mea-sured (>3), this dispersion was adopted whereas the direct measurement error was usedfor species probed by fewer lines.

Table 5.2: Adopted parameters of the stellar atmosphere model for each starStar ID Teff(K) Log g [Fe/H] vt(km/s)Cl1-D56 4600 1.0 -2.60 2.1Cl1-D68 4350 0.5 -2.60 2.0Cl1-D164 4400 0.8 -2.60 2.1Cl2-B71 4450 0.7 -2.10 1.8Cl2-B77 4350 0.7 -2.10 1.7Cl2-B226 4250 0.6 -2.10 2.0Cl3-B59 4400 0.5 -2.30 2.0Cl3-B61 4400 0.8 -2.30 1.8Cl3-B82 4350 0.5 -2.30 2.0M15-S1 4350 0.5 -2.40 1.9M15-S3 4400 0.6 -2.40 1.8M15-S4 4150 0.6 -2.30 2.3M15-S6 4400 0.7 -2.40 1.8M15-S7 4400 0.4 -2.50 1.9

Table 5.3: Dependencies on model atmosphere parameters∆ Teff ∆ Log g ∆ vt Combined-200 K -0.3 -0.2 km s−1

[Ba ii/Fe i] -0.11 0.11 -0.06 0.17[Ca i/Fe i] 0.00 -0.01 0.05 0.05[Cr i/Fe i] 0.21 -0.03 -0.02 0.21[Eu ii/Fe i] -0.16 0.13 0.08 0.22[Fe i/H] 0.27 -0.05 -0.10 0.29[Fe ii/H] -0.05 0.07 -0.06 0.10[La ii/Fe i] -0.12 0.12 0.08 0.19[Mg i/Fe i] -0.06 -0.08 0.02 0.10[Mn i/Fe i] -0.02 -0.01 0.08 0.08[Na i/Fe i] 0.18 -0.04 -0.04 0.19[Nd ii/Fe i] -0.14 0.10 0.04 0.18[Ni i/Fe i] 0.02 -0.02 -0.03 0.04[O i/Fe i] -0.11 0.13 0.08 0.19[Ti i/Fe i] 0.27 -0.02 0.02 0.27[Ti ii/Fe i] -0.21 0.09 0.02 0.23[Y ii/Fe i] -0.18 0.09 0.04 0.21[Zn i/Fe i] -0.28 0.03 0.06 0.29

5.4: Interpretation 85

However, there is more than just the measurement error to consider. The chosen stel-lar model will also affect the derived abundances. The three important parameters in themodel are: temperature, gravity and micro-turbulence velocity. Each of these influencesthe final abundance in a different way. We estimated the uncertainty in each of thesethree parameters using the corresponding statistical errors on the slopes and the (Fe i -Fe ii) difference and computed the resultant change in abundance for all elemental ratios.Table 5.3 shows the abundance offset generated by each parameter, and the combinedeffect of all three (added quadratically).

A summary of our abundance analysis is available in Table 5.A2 (Fornax) and Ta-ble 5.A3 (M15) where we present all of our elements with the associated error estimatesand the number of lines used to compute the ratio. Only the EW measurement error isused in plots and tables, and Fe i is used to determine our [el/Fe] ratios.

5.4 Interpretation5.4.1 The Iron abundanceTable 5.4 compares our mean [Fe/H] with the latest results of two different classical meth-ods: RGB slope fitting and integrated spectroscopy. Our abundances appear 0.3 to 0.5dex lower than previous estimates. Most of this discrepancy is attributable to differentreference calibrators. Indeed, both the integrated spectroscopy and isochrone fitting arebased on the Zinn & West (1984) metallicity scale which places M15 at 〈[Fe/H]〉= −2.15and M92 at 〈[Fe/H]〉= −2.24. In contrast, high resolution spectroscopic analyses consis-tently find 〈[Fe/H]〉=−2.4 for M15, including Sneden et al. (1997) and this present work.Meanwhile, M92 is found to be 〈[Fe/H]〉=−2.34 (Sneden et al., 2000). The differencebetween high resolution spectroscopy and the other indirect methods, due to differencesin calibration, is therefore of the order of 0.25 dex, the rest of the discrepancy might bedue to the propagation of errors, and indeed appears of the order of the quoted errorbars (± 0.2dex). In conclusion, although the absolute value of metallicities presented inthe works quoted in Table 5.4 do not appear accurate, the comparison made by the au-thors with M15 and M92, the most metal-poor clusters known in our Galaxy, is correct.Our analysis reveals that Cluster 1, at 〈[Fe/H]〉= −2.5, is actually the most metal-poorglobular cluster yet observed. It is clearly more metal-poor than M15, with weaker ironlines, as can be seen in Figure 5.5, where we compare Cl1-D68 and M15-S1, two RGBstars of similar temperature, surface gravity and micro-turbulence velocity.

Table 5.4: Recent metallicity estimates from different methods

Cluster 1 Cluster 2 Cluster 3 Method Reference−2.5± 0.1 −2.1± 0.1 −2.4± 0.1 Individual stars, HR spectra This work

N/A −1.76± 0.41 −1.84± 0.18 Integrated light spectra Strader et al. 2003−2.20± 0.20 −1.78± 0.20 −1.96± 0.20 RGB Slope Buonanno et al. 1998


Figure 5.5: The comparison between Cl1-D68 and M15-S1. These are two RGB starswith similar stellar parameters but a difference in [Fe/H] of 0.2 dex.

5.4.2 The Alpha elements

Alpha elements come predominantly from Type II supernovae, unlike Fe which comespredominantly from type Ia SN (McWilliam 1997, Tinsley 1979). The [α/Fe] ratiosfrequently display different patterns with respect to Fe in different environments (e.g.,Shetrone et al. 2001). They are typically overabundant by +0.3 to +0.4 dex in Galacticglobular cluster stars and halo stars with respect to solar, as expected in old componentswhere only SNe II have had time to contribute to the chemical enrichment.


Figure 5.6: Alpha elements abundances as a function of [Fe/H]. Filled triangles are forCluster 1, filled diamonds are for Cluster 2 and filled squares are for Cluster 3. Circlesare for our M15 stars. Small grey dots are galactic stars and small empty circles aregalactic GCs. Upper limits, when present, are shown with one sided arrows, replacingthe error bars. See text for more details.

In Figure 5.6, we plot the abundance ratios for α-elements Ca, Mg and O in FornaxdSph globular clusters 1, 2 & 3. Also plotted are the four M15 control stars and, assmaller symbols, Galactic halo stars, taken from the compilation of Venn et al. (2004)and Galactic globular cluster stars from the compilation (averaged by cluster) of Pritzlet al. (2005), except for [O/Fe] points, which are from Shetrone et al. 1996a (individualstars, not averages.) The abundances of Ca, Mg and O are all above the solar value (justlike the Galactic halo and globular cluster stars) with a small dispersion and small error


Figure 5.7: Titanium abundances as a function of [Fe/H]. The symbols are the same asin Figure 5.6. Separated Ti i and Ti ii were not available for our halo stars, so a global[Ti/Fe] is used for these points.

bars. There are a couple of Fornax dSph globular cluster stars with clearly anomalous Oand Mg abundance, and they will be discussed later in section 5.4.3. The Fornax dSphglobular cluster α-element ratios appear to follow the same patterns found in Galac-tic globular cluster stars, suggesting that the oldest epoch of globular cluster formationis very similar in these two different environments. The overabundance of α-elementsseen in Galactic globular clusters stars may be interpreted as the number of massivestars present in the early history of our Galaxy assuming that the main contributor toα-elements is SNe II explosions from massive stars. The same over abundance is seenin Fornax dSph globular clusters so this enrichment pattern is not only present in ourGalaxy.

Titanium is shown in Figure 5.7, where we chose to compare our results with theGalactic globular clusters studied by Shetrone et al. (2003) rather than the compilationof Pritzl et al. (2005) for homogeneity purposes. At first glance, Ti seems to be under-abundant in the Fornax clusters with respect to halo stars in the Milky Way. However,they fall right on top of our M15 and Shetrone’s M30, M68 and M55, close to a solarTi/Fe ratio. However, we would like to stress that [Ti/Fe] ratios of different authors canbe on different scales (depending on the set of Ti lines used and the adopted log gfs), aswell illustrated by M15: the [Ti/Fe] ratios in M15 found by Sneden et al. (1997, includedin the Pritzl compilation) are ∼0.4 dex higher than in our own analysis of M15, but using


Figure 5.8: Here we show the “Deep-Mixing” abundance anomaly. An anti-correlationof O-Na on the left and a correlation of O-Mg on the right. The symbols are the sameas in Figure 5.6.

Sneden’s Ti lines, log gfs and EWs (in the stars we have in common), our analysis yieldsthe same value as Sneden’s. We also notice a small systematic difference between theratio of Ti i and Ti ii over iron (∼0.2 dex), that could be caused by log gfs (that couldbe on different scales for Ti i and Ti ii) and/or non-LTE effects . We therefore concludethat, based on the comparison of our Fornax globular clusters with a fully compatibleanalysis of galactic globular clusters (our analysis of M15 and three other clusters byShetrone et al. 2003), there is no difference in the Ti/Fe ratios observed in Fornax andMW globular clusters.

5.4.3 Deep mixing pattern

Deep-mixing occurs when material processed deep inside a star finds its way to the upperatmosphere, thus modifying the original abundance pattern. Proton-capture nucleosyn-thesis converts O, N, Ne to Na, and Mg to Al in the H fusion layer of evolved RGBstars. This means that a significant atmospheric depletion of O caused by deep-mixingshould be accompanied by an enhancements of Na (Langer et al. 1993) and similarlyan enhancement in Al should cause observable Mg depletion (Langer & Hoffman 1995).Such patterns (anti-correlations of O-Na and Mg-Al) are found in galactic globular clus-ter stars but not in comparable field stars of our Galaxy (Gratton et al. 2004), or anyother (e.g., Shetrone et al. 2001). It is assumed that this is caused by environmentaleffects within a star cluster but whether it is the result of deep-mixing within the RGBstars that are observed or the fossil traces of self-pollution of the globular cluster duringits formation process, or a combination of the two, is not well understood.


Figure 5.8 shows that deep mixing patterns are not only found in galactic globularclusters and the old clusters of the Large Magellanic Cloud [LMC] (Hill et al. 2000)but also in clusters of much smaller dwarf spheroidal galaxies like Fornax. The anti-correlation O-Na is visible in two (Cl3-B82 and Cl1-D164) of the nine stars we observedin the Fornax globular clusters displaying high Na and low O abundances (left panel),accompanied by low Mg abundances (correlation O-Mg, right panel). We cannot checkdirectly whether the Mg-Al anti-correlation also exists in these clusters, since we did notdetect Al in our Fornax dSph globular cluster spectra, because the Al lines present in ourspectral range are too weak. Our detection limit is about 14 mÅ, which translates into anupper limit to [Al/Fe] of 1.4. Shetrone et al. (1996b) found that the usual enhancementof Al ranges from 0.5 to 1.0 dex, thus largely consistent with our upper limit.

5.4.4 Iron-peak elementsThe Fe-peak elements we observed in the Fornax dSph globular clusters are Cr, Ni andZn, and they are shown in Figure 5.9. Comparison points for Galactic globular clustersare from the compilation of Pritzl et al. (2005), and Galactic halo stars are from theHamburg-ESO (HERES) survey (Barklem et al. 2005) for Cr (top panel), from the com-pilation of Venn et al. (2004) for Ni (middle panel), and from Sneden et al.(1991) andBarklem et al. (2005) for Zn (lower panel).

Cr is believed to be produced mainly by incomplete explosive silicon burning (Woosley& Weaver 1995). Despite large error bars in our measurements of the Fornax dSph glob-ular cluster stars, there seems to be an increase (by ∼0.3 dex) of the [Cr/Fe] ratiobetween the two more metal-poor clusters and the more metal-rich Cluster 2. Such atrend of increasing [Cr/Fe] with increasing [Fe/H] has been observed in Galactic fieldstars (McWilliam et al. 1995, Carretta et al. 2002), leading to a similar ∼0.3 dexincrease, but over a much wider metallicity range (−3.5 to −2.). Newer, high qualityobservations by Cayrel et al. (2004) of Galactic halo stars further reduced the observedslope of increasing [Cr/Fe] with increasing metallicity to ∼0.15 dex over a [Fe/H] rangefrom −2.5 to −4 dex, with an extremely small intrinsic scatter (σ = 0.05 dex). Thehigher [Cr/Fe] abundance observed in Cluster 2 therefore seems unlikely, and is probablycaused by our observational errors.

Ni is believed to be produced in complete explosive silicon burning. We don’t expectany relation between [Ni/Fe] as a function of [Fe/H], based on what we see in the MW.Even at this low metallicity, the relation is flat with a value close to zero, within theerror bars, as we can see in Figure 5.9. This is consistent with the majority of Galacticglobular clusters, open clusters and halo stars (Sneden et al. 2004). So yet again, theFornax dSph globular clusters are similar to the normal Galactic globular clusters.

Zn has the same origin as Ni, but it has been suggested (Heger & Woosley 2002)that it could also be formed by neutron capture, and be either an r-process or an s-process element. Our results, more than half of which are upper limits, are consistentwith Galactic values.


Figure 5.9: Iron-peak elements abundances as a function of [Fe/H]. The symbols arethe same as in Figure 5.6.


Figure 5.10: Heavy elements abundances as a function of [Fe/H]. The symbols are thesame as in Figure 5.6.

5.4.5 Heavy elements

The heavy elements Y, Ba and Eu in the Fornax dSph globular clusters are plotted inFigure 5.10. [Y/Fe] appears to be consistent with what is observed in Galactic globularclusters. However, Cluster 1 and 3 (the two most metal-poor) appear to have higher[Ba/Fe] than average for Galactic globular clusters. As shown in Figure 5.4), europiumis measured from a single weak line, and could only be detected in Cluster 3 (all otherFornax points in this plot are upper limits), in which [Eu/Fe] is particularly high, abovethe typical range for Galactic globular clusters.


Ba and Y are neutron-capture elements which are, in the solar system, dominatedby the s-process, a process due to low to intermediate-mass Asymptotic Giant Branch(AGB) stars, with only a minor contribution from the r-process. Eu on the other hand isalmost entirely dominated by the r-process, which requires more extreme neutron fluxes,such as SN II explosions (McWilliam 1997) associated with massive stars. In the MilkyWay, with decreasing metallicities the s-process contribution gradually decreases (con-sistent with the timescale of AGB evolution) so that below ∼ −2.5dex, both in field andglobular clusters stars, all heavy elements are dominated exclusively by the r-process.(Johnson et al. 2001, James et al. 2004, Barklem et al. 2005). In Cluster 3, we detect,not only Eu, but also other heavy elements represented by weak lines preventing detec-tion in the other clusters: Nd and La.

In Figure 5.11, we compare Cluster 3 log (ε)∗ values to the solar system r− and s−process abundances (Burris et al. 2000). The solar system elemental abundances havebeen shifted by the difference between the mean values of Eu for Cluster 3 and the solarsystem abundance distribution (−1.55 dex). Clearly, the abundances of most elements inthe Fornax globular clusters match the solar system r-process pattern within the obser-vational uncertainties (with the exception of La which seems to be matched by neitherthe r− nor the s− process patterns). Cluster 3 stars are obviously very close to ther-process expectations, confirming that, similarly to the most metal-poor globular clus-ters in the galactic Halo (M15, M92, M68), cluster 3 is also dominated by the r-process.This is also confirmed by the [Ba/Eu] ratio observed in the three Cluster 3 stars [Ba/Eu]= −0.62, −0.69, −0.7 (±0.20), very close to the −0.69 for the r-process component inthe solar system (as compared to +1.15 for the s-process, Arlandini et al. 1999). Theupper limits for Eu in Clusters 1 and 2, although not decisive, are also compatible witha pure r-process enrichment ([Ba/Eu]> −1.00 to −0.82). This result indicates that, inFornax dSph as in our Galaxy, heavier neutron capture elements in the lowest metallicitystars have only very weak s-process contribution. Or in other words, that heavy elementsin Fornax dSph globular clusters, as in M15, are formed principally through the r-process.

The high neutron-capture element content of Cluster 3, that we attribute to the r-process, is similar for all three stars, and above the upper edge of the range of valuestraditionally covered by the Galactic globular clusters (Pritzl et al. 2005). R-processenrichments of this order or even higher are found in Galactic halo field stars (Barklemet al. 2005), but as far as Galactic globular clusters are concerned, the only case knownto date is M15. Sneden et al. (1997, 2000) have established that M15 has a stellarbi-modality with one group being strongly overabundant in [Eu/Fe] (and [Ba/Fe]) com-pared to the other. Our observations of three stars in Cluster 3 do not provide sufficientstatistics to determine if this cluster also has a bimodality (with our 3 stars by chancehappening to belong to the high Eu group) or if all stars in Cluster 3 are Eu-rich.

Finally, despite the dispersion in Ba that seems to exist among the three Fornaxglobular clusters (Cluster 3 being the most Ba and Eu rich), Y is very similar fromcluster to cluster, and comparable to the Galactic abundances of this element (globularclusters and field stars). This also leads Cluster 3 to have a Ba/Y ratio higher than inthe two other clusters ([Ba/Y]=+0.43 compared to +0.0 in Cluster 2, and marginally∗ The scale used for log (ε) is the standard astronomical scale (log10(Nel/NH) + 12).


Figure 5.11: The relative contributions of the r− and s− processes for the heavyelements in Cluster 3 (filled circle). The solar r− and s− process abundances, traced bya dotted and a full line respectively, are taken from Burris et al. (2000). They are shiftedby the difference between Cluster 3 and the solar system abundance for Eu (−1.55 dex).

higher than the +0.28 dex observed in Cluster 1). Interestingly, the [Ba/Y] observedin the three Fornax clusters are yet again very similar to that of the Galactic globularclusters and halo field stars, whereas the (on average more metal-rich) field stars in dwarfspheroidal galaxies have been shown to display systematically higher [Ba/Y] than theirgalactic counterparts (Venn et al. 2004).

5.5: Conclusions 95

Figure 5.12: The cluster mean elemental abundances of the three Fornax dSph globularclusters and M15. Each individual stellar abundance has being weighted by its error.Cluster 1 is identified by a filled square, Cluster 2 by a star, Cluster 3 by a cross andM15 by a triangle.

5.5 ConclusionsWe have compared the properties of the globular clusters belonging to the Fornax dSphwith those of the Milky Way with unprecedented accuracy. The Fornax dSph containsclusters with a range of properties such as metallicity, central concentration and Horizon-tal Branch structure. For the first time detailed chemical abundances have been derivedfor a sample of stars in a globular cluster system in an external galaxy, apart from theMagellanic Clouds. Despite their very different mass, morphology and global star forma-tion history, the Fornax dSph and the Milky Way appear to have experienced the same


very early enrichment conditions and in particular similar nucleosynthesis. This is sum-marised in Figure 5.12, where the mean elemental abundances, each being weighted by itserror, of the three Fornax globular clusters and M15 are compared. The abundance pat-terns of the individual stars in Milky Way globular clusters and Fornax globular clustersmatch each other almost perfectly. We find that the star-to-star abundance dispersionin the Fornax clusters is modest and compatible with similar observations of Galacticglobular clusters.

We have definitively established that the Fornax globular Clusters 1, 2 and 3 are verymetal-poor, slightly poorer than previous estimates, with respectively 〈[Fe/H]〉=−2.5,−2.1 and −2.4. Part of the discrepancy with previous studies is explained by the differ-ent reference calibrations used. Cluster 1 is now the most metal-poor globular clusterknown, however the difference between Cluster 1 and M92 or M15 in the Milky Way issmall. There seems to be universal lower limit to the metallicity at which star clustersform, which is higher than that of field stars in the halo of our Galaxy, where significantnumbers of stars are found with [Fe/H] < −4. It is also clear, that as in our Galaxy, theratio of the number of globular cluster to the number of field stars strongly decreaseswith rising metallicity (Harris & Harris 2002).

Clusters 1, 2 and 3 were clearly formed promptly and early in the history of FornaxdSph, alike the Milky Way globular clusters. They are over abundant in α-elements (O,Mg, Ca) at a similar level to Galactic clusters at identical [Fe/H], and the heavy elementabundances (Y, Ba, Eu) in the 3 clusters are compatible with dominant r-process en-richment. Finally, the Fe-peak elements are also very similar to Galactic globular clustervalues, with [Ni/Fe] being unambiguously solar in all three clusters and Zn and Cr arealso compatible with Galactic values.

The analogy between Galactic and Fornax dSph holds even in the rare cases andanomalies: (i) Eu is extremely overabundant in Cluster 3 stars. The only Galactic coun-terpart known to date is M15. (ii) Cl1-D164 and Cl3-B82 show low O and Mg associatedwith a high Na abundance, thus establishing an O-Na anti-correlation and O-Mg corre-lation. This is the same deep-mixing pattern observed in Galactic star clusters, and oldLMC clusters.

The effort towards a comprehensive description of the formation and evolution ofthe Fornax dSph will soon benefit from the analysis of VLT/FLAMES high resolutionspectra of a hundred field stars (Letarte et al., in preparation). It will then be possibleto describe the chemical enrichment and nucleosynthetic processes dominant for the fieldstar population compared to that found in the globular clusters, and to see when and ifthe similarities in enrichment patterns with our Galaxy end.


Acknowledgements:We gratefully acknowledge Carlo Emanuele Corsi and Roberto Buonanno for providingtheir photometry for our target selection. ET gratefully acknowledges support from afellowship of the Royal Netherlands Academy of Arts and Sciences. BL is funded by agrant from the Netherlands Organisation for Scientific Research (NWO). PJ and GMgratefully acknowledge support from the Swiss National Science Foundation (SNSF).


Table 5.A1: Line parameters and equivalent widths for the Fornax globular clusters and M15.When there is a * in the S column, it indicates that a synthetic spectra was used for the abundancedetermination. HFS indicates a line with hyperfine splitting, so no individual EW measurement forthat line is available. (Part 1)

λ El χex log gf S Cl1-D56 Cl1-D68 Cl1-D164 Cl2-B71 Cl2-B77 Cl2-B2264934.12 Ba ii 0.00 -0.703 * HFS HFS HFS HFS ... ...5853.69 Ba ii 0.60 -1.010 * HFS HFS HFS HFS HFS HFS6141.73 Ba ii 0.70 -0.077 * HFS HFS HFS HFS HFS HFS6496.91 Ba ii 0.60 -0.380 * 103.7 136.8 123.6 121.5 115.0 138.26102.73 Ca i 1.88 -0.790 50.4 54.5 50.8 91.9 97.3 116.06122.23 Ca i 1.89 -0.320 85.0 93.5 86.5 129.0 119.2 160.16161.30 Ca i 2.52 -1.270 ... ... ... ... ... 26.36166.44 Ca i 2.52 -1.140 ... ... ... 27.7 ... 32.86169.04 Ca i 2.52 -0.800 30.2 14.6 12.7 33.3 41.7 51.56169.56 Ca i 2.52 -0.480 ... 24.1 17.9 41.1 45.7 58.86439.08 Ca i 2.52 0.390 77.5 69.2 72.4 103.1 114.0 139.86455.60 Ca i 2.52 -1.290 ... ... ... ... ... 30.36499.65 Ca i 2.52 -0.820 24.1 22.1 11.2 31.3 35.1 58.95206.04 Cr i 0.94 0.019 100.0 124.2 109.0 126.4 170.5 214.45409.80 Cr i 1.03 -0.720 66.3 74.2 74.8 95.5 114.1 148.96645.13 Eu ii 1.37 0.200 * ... HFS HFS HFS HFS HFS4966.10 Fe i 3.33 -0.890 51.3 54.0 39.8 68.2 84.0 85.15006.12 Fe i 2.83 -0.628 82.0 91.4 108.2 117.1 112.4 140.05079.75 Fe i 0.99 -3.240 ... 127.1 119.6 141.9 136.6 193.05083.35 Fe i 0.96 -2.862 100.3 124.7 122.3 138.8 135.0 181.75150.85 Fe i 0.99 -3.030 84.7 114.4 112.7 117.4 129.9 178.35151.92 Fe i 1.01 -3.326 69.7 115.4 90.3 104.4 118.9 155.95162.29 Fe i 4.18 0.020 64.5 32.0 44.7 76.3 67.2 90.25166.28 Fe i 0.00 -4.200 110.5 ... 130.8 140.5 180.3 198.85171.61 Fe i 1.48 -1.751 126.3 133.7 126.4 132.6 149.0 178.15192.34 Fe i 3.00 -0.520 98.0 89.1 79.0 104.7 124.7 134.15196.08 Fe i 4.26 -0.450 ... ... ... 41.8 9.3 33.25215.19 Fe i 3.27 -0.930 41.4 35.8 43.9 76.4 81.9 86.75216.28 Fe i 1.61 -2.102 97.6 104.9 107.0 115.1 128.2 165.45217.30 Fe i 3.21 -1.270 36.0 ... 45.4 63.4 ... 93.35232.95 Fe i 2.94 -0.067 122.3 122.8 104.6 130.9 135.9 166.25250.21 Fe i 0.12 -4.700 59.0 78.4 78.8 95.0 99.2 145.1



λ El χex log gf S Cl1-D56 Cl1-D68 Cl1-D164 Cl2-B71 Cl2-B77 Cl2-B2265307.37 Fe i 1.61 -2.812 58.2 76.3 70.9 91.7 92.4 129.15324.19 Fe i 3.21 -0.100 86.4 84.0 100.6 127.7 114.8 146.25339.93 Fe i 3.27 -0.680 65.8 59.3 52.1 111.3 99.2 123.65364.86 Fe i 4.45 0.220 ... ... 32.8 52.1 92.8 72.25367.48 Fe i 4.42 0.550 34.3 ... 37.7 54.7 56.7 78.25369.96 Fe i 4.37 0.540 64.2 ... 26.0 74.3 68.3 88.65371.50 Fe i 0.96 -1.644 205.6 195.4 188.1 171.9 196.1 273.15383.37 Fe i 4.31 0.500 49.7 49.1 51.4 79.1 59.9 102.95393.17 Fe i 3.24 -0.920 58.3 57.2 70.4 80.9 116.0 111.75397.14 Fe i 0.91 -1.992 161.7 181.1 172.3 182.3 198.0 246.35405.79 Fe i 0.99 -1.852 159.2 190.7 163.8 173.6 201.6 230.35415.19 Fe i 4.39 0.510 54.3 ... 50.3 68.3 91.9 96.05424.07 Fe i 4.32 0.520 72.8 67.7 57.2 80.9 94.1 91.25501.48 Fe i 0.96 -3.050 114.1 132.9 124.7 143.4 140.8 174.05506.79 Fe i 0.99 -2.790 111.2 131.9 134.3 157.9 149.1 206.15615.66 Fe i 3.33 0.050 84.7 88.8 94.9 115.0 124.5 149.45956.70 Fe i 0.86 -4.570 40.8 37.5 ... 51.1 ... 99.26024.05 Fe i 4.55 -0.110 18.6 12.0 ... 28.8 49.1 54.36136.62 Fe i 2.45 -1.500 98.8 91.3 92.3 119.8 109.6 159.36137.70 Fe i 2.59 -1.366 77.7 77.1 83.7 109.1 113.7 143.96157.75 Fe i 4.07 -1.260 ... ... ... ... ... ...6173.34 Fe i 2.22 -2.850 29.8 21.5 21.8 45.8 62.2 80.56191.57 Fe i 2.43 -1.416 76.3 92.6 93.4 121.7 117.8 140.86213.43 Fe i 2.22 -2.660 34.6 39.9 41.0 68.5 59.0 91.06219.29 Fe i 2.20 -2.438 46.6 60.3 52.4 87.5 86.6 109.06229.23 Fe i 2.84 -2.900 ... ... ... ... 33.0 24.76230.74 Fe i 2.56 -1.276 84.1 93.2 94.6 115.6 49.3 156.86232.64 Fe i 3.65 -0.960 ... ... 16.5 40.3 45.1 45.36252.57 Fe i 2.40 -1.757 82.8 82.9 98.1 110.2 124.0 146.56270.23 Fe i 2.85 -2.610 ... ... ... 16.9 24.1 37.56297.80 Fe i 2.22 -2.740 ... 57.3 61.8 68.8 80.5 113.56301.50 Fe i 3.65 -0.720 62.0 51.3 44.4 68.0 81.4 106.56302.49 Fe i 3.69 -1.150 ... 9.2 8.8 23.8 51.0 57.16393.61 Fe i 2.43 -1.630 68.8 94.3 95.9 86.4 137.7 154.76421.36 Fe i 2.28 -2.014 52.9 72.9 71.2 90.4 136.0 142.06430.86 Fe i 2.18 -1.946 66.9 87.8 88.5 94.7 ... 148.36481.87 Fe i 2.27 -2.980 ... 27.9 23.9 42.8 57.0 79.86498.94 Fe i 0.96 -4.690 ... ... 30.1 38.1 57.7 98.76518.37 Fe i 2.83 -2.460 ... ... ... ... 23.8 47.76574.23 Fe i 0.99 -5.020 ... 11.5 7.5 ... 38.2 64.36593.88 Fe i 2.43 -2.390 ... 33.6 27.2 69.8 66.1 89.16609.12 Fe i 2.56 -2.660 ... 15.1 6.7 37.4 55.7 63.64923.92 Fe ii 2.89 -1.320 117.7 125.4 122.9 130.4 114.8 132.45197.57 Fe ii 3.23 -2.100 58.3 43.8 49.7 65.6 77.3 78.45234.63 Fe ii 3.22 -2.118 64.1 52.6 52.1 ... 69.2 80.05276.00 Fe ii 3.20 -1.950 96.5 50.1 55.3 87.9 93.2 92.45284.10 Fe ii 2.89 -3.190 ... ... 21.2 ... 37.0 52.85325.56 Fe ii 3.22 -2.600 ... ... ... ... ... ...5425.25 Fe ii 3.20 -3.360 ... ... 22.8 34.5 ... ...5534.85 Fe ii 3.24 -2.920 27.9 27.8 24.4 46.8 42.7 43.9



λ El χex log gf S Cl1-D56 Cl1-D68 Cl1-D164 Cl2-B71 Cl2-B77 Cl2-B2265991.38 Fe ii 3.15 -3.740 ... ... ... ... ... ...6149.25 Fe ii 3.89 -2.720 ... ... ... ... 24.1 ...6238.38 Fe ii 3.89 -2.480 ... ... ... ... ... ...6369.46 Fe ii 2.89 -4.250 ... ... ... ... ... ...6432.68 Fe ii 2.89 -3.710 ... 17.5 ... 19.0 ... ...6456.39 Fe ii 3.90 -2.080 ... 6.8 19.8 ... ... 32.26516.08 Fe ii 2.89 -3.450 ... ... ... 35.9 35.1 40.85301.97 La ii 0.40 -1.140 * ... ... ... ... ... ...5303.52 La ii 0.32 -1.350 * ... ... ... ... ... ...6320.43 La ii 0.17 -1.562 * ... ... ... ... ... ...6390.46 La ii 0.32 -1.400 * ... ... ... ... ... ...6774.27 La ii 0.13 -1.708 * ... ... ... ... ... ...5172.70 Mg i 2.71 -0.390 * 283.0 266.1 241.8 312.6 340.9 ...5528.41 Mg i 4.35 -0.357 * 110.2 93.3 79.0 130.7 141.1 135.85711.09 Mg i 4.35 -1.728 * 26.2 12.1 8.1 48.5 46.9 53.16013.51 Mn i 3.07 -0.252 * ... ... ... ... ... 25.96021.82 Mn i 3.08 0.035 * ... ... ... ... ... 40.95889.97 Na i 0.00 0.122 * 230.5 219.7 249.6 230.4 231.4 ...5895.94 Na i 0.00 -0.184 * 204.0 185.1 240.0 216.1 215.6 ...6154.23 Na i 2.10 -1.560 ... ... ... ... ... ...5249.59 Nd ii 0.98 0.217 * ... 27.5 ... ... ... 21.15319.82 Nd ii 0.55 -0.194 * 37.9 30.8 30.7 ... ... 33.55476.92 Ni i 1.83 -0.890 152.1 124.5 96.3 96.5 131.6 158.16176.82 Ni i 4.09 -0.430 ... ... 8.6 22.1 ... ...6177.25 Ni i 1.83 -3.500 ... ... 3.4 ... ... 21.16300.31 O i 0.00 -9.760 * 19.3 11.9 10.0 21.6 23.8 45.04840.87 Ti i 0.90 -0.450 40.6 23.4 17.6 48.8 64.8 69.84913.62 Ti i 1.87 0.216 ... ... 18.5 39.0 ... 41.25014.24 Ti i 0.81 0.910 97.2 96.3 95.3 154.4 173.8 204.85016.16 Ti i 0.85 -0.510 31.4 26.5 ... 43.7 69.2 80.35064.65 Ti i 0.05 -0.930 60.9 73.3 63.8 86.1 99.7 148.05210.39 Ti i 0.05 -0.580 60.5 76.3 77.2 98.1 115.0 164.94798.53 Ti ii 1.08 -2.670 54.1 36.3 ... 45.6 ... ...5129.16 Ti ii 1.89 -1.390 63.2 54.4 55.5 63.4 104.1 96.55154.07 Ti ii 1.57 -1.520 62.2 69.6 71.2 ... 78.7 96.15226.55 Ti ii 1.57 -1.000 90.2 86.7 97.6 105.2 124.6 126.25381.01 Ti ii 1.57 -1.780 54.2 39.9 47.2 71.8 81.5 97.15418.77 Ti ii 1.58 -2.110 ... 40.6 55.5 67.4 60.5 67.94883.69 Y ii 1.08 0.070 41.4 41.0 37.2 56.8 82.3 73.44900.11 Y ii 1.03 -0.090 80.2 ... ... 65.6 ... 160.85087.43 Y ii 1.08 -0.170 ... 32.9 29.2 66.4 38.9 64.05200.42 Y ii 0.99 -0.570 34.5 29.9 ... 49.4 40.9 53.54810.54 Zn i 4.08 -0.170 * 22.0 40.9 31.9 32.9 43.0 50.0


Table 5.A1: Line parameters and equivalent widths for the Fornax globular clusters and M15.When there is a * in the S column, it indicates that a synthetic spectra was used for the abundancedetermination. HFS indicates a line with hyperfine splitting, so no individual EW measurement forthat line is available. (Part 2)

λ El χex log gf S Cl3-B59 Cl3-B61 Cl3-B82 M15S1 M15S3 M15S6 M15S74934.12 Ba ii 0.00 -0.703 * HFS ... ... HFS HFS HFS HFS5853.69 Ba ii 0.60 -1.010 * ... HFS ... HFS HFS HFS HFS6141.73 Ba ii 0.70 -0.077 * HFS HFS HFS HFS HFS HFS HFS6496.91 Ba ii 0.60 -0.380 * 144.8 123.1 147.2 124.8 136.2 149.2 118.46102.73 Ca i 1.88 -0.790 60.2 63.4 72.4 77.7 74.8 79.9 71.56122.23 Ca i 1.89 -0.320 107.0 104.5 103.8 116.1 110.6 114.2 105.96161.30 Ca i 2.52 -1.270 ... ... ... 11.8 14.0 9.7 9.06166.44 Ca i 2.52 -1.140 ... ... ... 17.6 16.7 ... 16.96169.04 Ca i 2.52 -0.800 ... 21.4 33.6 28.4 28.4 32.0 24.76169.56 Ca i 2.52 -0.480 ... 41.0 ... 43.4 39.2 44.2 38.36439.08 Ca i 2.52 0.390 84.5 88.2 97.1 99.7 95.7 99.4 95.96455.60 Ca i 2.52 -1.290 26.0 ... ... 13.1 10.5 9.5 10.06499.65 Ca i 2.52 -0.820 ... 20.0 25.6 26.8 23.4 26.4 23.65206.04 Cr i 0.94 0.019 136.3 129.2 129.3 126.9 115.7 121.3 112.25409.80 Cr i 1.03 -0.720 89.3 85.3 86.3 88.1 78.8 84.5 74.06645.13 Eu ii 1.37 0.200 * HFS HFS HFS HFS HFS HFS HFS4966.10 Fe i 3.33 -0.890 53.5 42.6 57.8 63.3 60.4 64.2 51.95006.12 Fe i 2.83 -0.628 119.5 109.7 113.7 107.0 100.9 106.7 97.35079.75 Fe i 0.99 -3.240 133.1 116.6 130.0 118.6 109.7 115.2 101.75083.35 Fe i 0.96 -2.862 126.3 114.5 134.0 130.5 120.6 122.2 114.85150.85 Fe i 0.99 -3.030 150.6 114.2 128.0 122.4 106.3 111.9 100.65151.92 Fe i 1.01 -3.326 114.1 120.3 143.8 111.5 ... 102.6 91.45162.29 Fe i 4.18 0.020 50.0 62.3 46.7 49.9 44.1 49.9 43.05166.28 Fe i 0.00 -4.200 153.3 159.4 146.4 144.6 128.0 131.9 121.05171.61 Fe i 1.48 -1.751 141.7 142.0 142.3 145.8 133.8 142.0 133.75192.34 Fe i 3.00 -0.520 122.7 82.0 110.5 102.2 96.8 97.4 96.05196.08 Fe i 4.26 -0.450 18.7 23.4 ... 11.8 9.4 8.8 9.15215.19 Fe i 3.27 -0.930 58.2 47.0 47.3 53.8 52.7 56.0 47.65216.28 Fe i 1.61 -2.102 130.4 119.2 135.9 123.1 110.6 118.4 108.55217.30 Fe i 3.21 -1.270 44.5 43.2 ... 50.5 48.7 49.9 ...5232.95 Fe i 2.94 -0.067 125.1 113.3 123.2 124.5 118.8 122.1 112.35250.21 Fe i 0.12 -4.700 102.3 87.7 95.2 92.5 76.5 83.3 67.95307.37 Fe i 1.61 -2.812 74.3 74.1 83.6 77.8 68.2 73.3 65.85324.19 Fe i 3.21 -0.100 106.4 100.7 105.9 105.7 99.6 104.3 84.35339.93 Fe i 3.27 -0.680 86.6 71.5 92.7 72.7 66.7 72.1 59.05364.86 Fe i 4.45 0.220 53.6 45.5 40.7 40.5 37.6 40.5 32.35367.48 Fe i 4.42 0.550 51.5 47.8 42.6 45.8 42.8 46.5 38.85369.96 Fe i 4.37 0.540 72.1 55.3 59.0 49.9 50.6 50.5 47.55371.50 Fe i 0.96 -1.644 191.1 184.0 201.7 187.2 177.4 180.7 171.35383.37 Fe i 4.31 0.500 60.3 60.0 64.6 61.6 57.7 59.6 50.95393.17 Fe i 3.24 -0.920 76.1 70.8 65.3 71.6 63.6 69.0 61.25397.14 Fe i 0.91 -1.992 171.5 180.6 190.1 177.9 163.5 169.4 157.85405.79 Fe i 0.99 -1.852 183.9 184.4 182.0 175.3 165.0 172.6 157.75415.19 Fe i 4.39 0.510 48.4 41.0 58.6 59.9 55.5 57.1 50.4



λ El χex log gf S Cl3-B59 Cl3-B61 Cl3-B82 M15S1 M15S3 M15S6 M15S75424.07 Fe i 4.32 0.520 78.3 72.4 71.0 65.1 63.0 66.0 55.55501.48 Fe i 0.96 -3.050 133.0 137.9 149.6 135.2 120.4 127.1 114.55506.79 Fe i 0.99 -2.790 142.3 122.2 150.1 144.3 131.0 135.3 127.85615.66 Fe i 3.33 0.050 104.0 95.0 111.9 105.8 98.5 102.8 95.05956.70 Fe i 0.86 -4.570 49.8 61.8 55.8 49.9 37.9 37.0 29.06024.05 Fe i 4.55 -0.110 32.1 29.1 22.7 25.3 20.7 24.6 19.86136.62 Fe i 2.45 -1.500 100.0 104.5 122.5 110.5 98.9 106.1 94.46137.70 Fe i 2.59 -1.366 104.0 88.9 105.5 98.8 93.0 99.0 89.46157.75 Fe i 4.07 -1.260 ... ... ... 8.8 8.5 9.6 6.36173.34 Fe i 2.22 -2.850 34.4 32.3 36.8 40.2 32.9 39.9 33.06191.57 Fe i 2.43 -1.416 123.1 86.8 109.4 109.0 100.7 101.4 94.76213.43 Fe i 2.22 -2.660 45.2 50.0 47.8 57.9 52.1 55.5 45.56219.29 Fe i 2.20 -2.438 78.6 64.7 60.6 66.6 59.5 68.7 53.06229.23 Fe i 2.84 -2.900 ... ... 16.7 8.3 7.7 11.5 7.06230.74 Fe i 2.56 -1.276 107.2 105.2 116.5 115.2 105.4 112.6 102.26232.64 Fe i 3.65 -0.960 ... 19.0 ... ... ... ... ...6252.57 Fe i 2.40 -1.757 105.2 100.7 100.9 101.7 95.0 94.3 86.86270.23 Fe i 2.85 -2.610 ... ... 20.4 ... ... ... ...6297.80 Fe i 2.22 -2.740 ... 70.0 59.9 69.0 65.7 ... ...6301.50 Fe i 3.65 -0.720 50.7 53.1 50.1 45.2 43.7 44.7 36.36302.49 Fe i 3.69 -1.150 ... ... ... ... 21.9 24.8 ...6393.61 Fe i 2.43 -1.630 94.8 72.0 102.6 104.9 93.2 99.3 91.16421.36 Fe i 2.28 -2.014 91.1 80.7 90.5 94.7 83.9 91.0 79.86430.86 Fe i 2.18 -1.946 98.4 97.5 110.6 105.3 92.1 98.5 90.36481.87 Fe i 2.27 -2.980 25.7 37.5 35.0 ... ... ... ...6498.94 Fe i 0.96 -4.690 35.8 37.0 40.1 41.0 28.6 34.3 23.36518.37 Fe i 2.83 -2.460 ... 19.2 17.1 28.7 29.4 ... ...6574.23 Fe i 0.99 -5.020 ... 19.9 29.2 23.2 16.5 20.8 13.96593.88 Fe i 2.43 -2.390 57.6 42.6 59.5 52.9 43.2 53.2 39.46609.12 Fe i 2.56 -2.660 ... 36.6 ... 24.7 20.4 20.8 16.34923.92 Fe ii 2.89 -1.320 117.3 115.6 117.7 122.4 121.7 124.8 118.65197.57 Fe ii 3.23 -2.100 38.6 52.9 45.4 54.9 51.6 57.5 51.25234.63 Fe ii 3.22 -2.118 55.1 121.1 ... 59.0 58.9 59.7 51.45276.00 Fe ii 3.20 -1.950 87.8 68.5 77.3 69.7 69.7 68.3 61.45284.10 Fe ii 2.89 -3.190 40.7 41.9 40.5 29.9 27.5 31.6 24.65325.56 Fe ii 3.22 -2.600 9.1 ... ... 14.9 14.2 15.4 14.25425.25 Fe ii 3.20 -3.360 ... ... ... 14.1 10.8 13.0 9.35534.85 Fe ii 3.24 -2.920 38.6 ... 27.8 28.0 27.7 26.2 24.05991.38 Fe ii 3.15 -3.740 ... ... ... 14.2 8.4 14.8 10.16149.25 Fe ii 3.89 -2.720 ... ... ... 6.5 4.7 8.8 8.66238.38 Fe ii 3.89 -2.480 ... ... ... 13.8 12.2 14.9 11.86369.46 Fe ii 2.89 -4.250 ... ... ... 10.1 ... ... ...6432.68 Fe ii 2.89 -3.710 ... ... 20.0 18.1 16.7 20.3 18.26456.39 Fe ii 3.90 -2.080 ... ... ... 26.6 28.0 28.6 26.56516.08 Fe ii 2.89 -3.450 ... ... 27.9 31.6 26.7 26.4 23.35301.97 La ii 0.40 -1.140 * 16.0 ... 22.3 ... ... ... ...5303.52 La ii 0.32 -1.350 * 12.9 ... 13.7 ... ... ... ...6320.43 La ii 0.17 -1.562 * 19.3 15.6 16.9 ... ... ... ...6390.46 La ii 0.32 -1.400 * 18.2 25.8 23.4 ... ... ... ...6774.27 La ii 0.13 -1.708 * 14.2 9.2 15.2 ... ... ... ...



λ El χex log gf S Cl3-B59 Cl3-B61 Cl3-B82 M15S1 M15S3 M15S6 M15S75172.70 Mg i 2.71 -0.390 * 276.4 260.0 204.2 264.1 266.4 237.1 247.95528.41 Mg i 4.35 -0.357 * 96.0 102.6 60.0 106.7 101.5 81.1 108.15711.09 Mg i 4.35 -1.728 * 18.1 27.6 8.6 22.5 20.0 11.1 19.66013.51 Mn i 3.07 -0.252 * ... 15.0 15.6 ... ... ... ...6021.82 Mn i 3.08 0.035 * ... 25.0 26.2 ... ... ... ...5889.97 Na i 0.00 0.122 * 242.1 206.2 251.6 240.6 296.2 302.3 237.75895.94 Na i 0.00 -0.184 * 217.3 199.4 271.0 216.2 257.3 261.2 194.36154.23 Na i 2.10 -1.560 ... ... ... ... ... ... ...5249.59 Nd ii 0.98 0.217 * 45.2 14.4 46.3 12.5 21.8 23.0 8.35319.82 Nd ii 0.55 -0.194 * 55.2 37.2 62.8 ... 31.8 38.0 ...5476.92 Ni i 1.83 -0.890 134.7 110.2 123.4 115.9 108.4 108.8 100.96176.82 Ni i 4.09 -0.430 ... ... ... 9.4 7.8 9.2 8.86177.25 Ni i 1.83 -3.500 ... ... ... 9.0 5.8 6.2 5.26300.31 O i 0.00 -9.760 * 20.8 22.7 11.9 10.7 ... 7.2 10.04840.87 Ti i 0.90 -0.450 33.9 ... 37.9 40.1 36.2 35.8 30.24913.62 Ti i 1.87 0.216 25.5 22.0 ... 14.4 5.7 15.3 10.95014.24 Ti i 0.81 0.910 122.3 122.0 124.6 ... ... ... ...5016.16 Ti i 0.85 -0.510 43.2 ... 36.8 45.7 36.3 41.5 37.45064.65 Ti i 0.05 -0.930 88.4 78.8 97.6 89.2 75.6 81.8 74.65210.39 Ti i 0.05 -0.580 111.1 105.7 98.6 94.2 83.5 84.3 77.64798.53 Ti ii 1.08 -2.670 54.6 45.1 ... ... ... ... ...5129.16 Ti ii 1.89 -1.390 79.4 74.1 63.7 74.5 72.1 73.1 68.75154.07 Ti ii 1.57 -1.520 80.2 92.2 81.2 77.3 73.2 75.3 66.65226.55 Ti ii 1.57 -1.000 109.7 91.4 103.6 106.2 101.1 104.5 98.25381.01 Ti ii 1.57 -1.780 78.0 67.5 65.4 64.2 56.8 62.5 54.55418.77 Ti ii 1.58 -2.110 50.9 54.8 42.4 54.5 52.4 51.8 47.94883.69 Y ii 1.08 0.070 58.9 55.0 56.3 53.1 58.9 64.9 47.14900.11 Y ii 1.03 -0.090 ... ... ... 64.1 77.2 81.4 53.85087.43 Y ii 1.08 -0.170 47.0 30.2 42.0 36.3 39.8 47.6 34.95200.42 Y ii 0.99 -0.570 ... ... 36.1 23.4 ... ... 20.04810.54 Zn i 4.08 -0.170 * 36.9 34.8 36.6 30.2 33.2 30.9 26.4


Table 5.A2: Fornax Globular Clusters elemental ratios

Cl1-D56 σ Nlines Cl1-D68 σ Nlines Cl1-D164 σ Nlines

[Ba ii/Fe i] -0.13 0.06 4 0.06 0.09 4 0.07 0.07 4[Ca i/Fe i] 0.27 0.09 5 0.18 0.07 5 0.09 0.05 4[Cr i/Fe i] -0.36 0.26 2 -0.35 0.12 2 -0.43 0.16 2

[Eu ii/Fe i] ... ... 0 <1.04 0.00 1 <0.89 0.00 1[Fe i/H] -2.40 0.03 40 -2.55 0.03 39 -2.59 0.03 45

[Fe ii/H] -2.43 0.06 4 -2.62 0.11 5 -2.56 0.06 7[La ii/Fe i] ... ... 0 ... ... 0 ... ... 0[Mg i/Fe i] 0.52 0.12 3 0.38 0.07 3 0.08 0.07 3[Mn i/Fe i] ... ... 0 ... ... 0 ... ... 0[Na i/Fe i] -0.10 0.12 2 -0.15 0.08 2 0.42 0.12 2

[Nd ii/Fe i] 0.60 0.20 1 0.40 0.07 2 0.49 0.10 1[Ni i/Fe i] ... ... 0 0.27 0.23 1 -0.20 0.21 1[O i/Fe i] <0.68 0.00 1 0.28 0.10 1 0.37 0.10 1[Ti i/Fe i] 0.07 0.12 5 -0.22 0.09 5 -0.09 0.14 5

[Ti ii/Fe i] 0.06 0.10 5 -0.02 0.06 6 0.18 0.08 5[Y ii/Fe i] -0.21 0.28 2 -0.29 0.14 3 -0.33 0.26 2[Zn i/Fe i] <-0.10 0.00 1 <0.45 0.00 1 <0.29 0.00 1

Cl2-B71 σ Nlines Cl2-B77 σ Nlines Cl2-B226 σ Nlines

[Ba ii/Fe i] -0.19 0.10 4 -0.12 0.10 3 -0.38 0.10 3[Ca i/Fe i] 0.27 0.07 7 0.16 0.05 6 0.21 0.04 9[Cr i/Fe i] -0.38 0.19 2 -0.10 0.16 2 -0.03 0.22 2

[Eu ii/Fe i] <0.63 0.00 1 <0.88 0.00 1 <0.60 0.00 1[Fe i/H] -2.14 0.03 50 -2.09 0.04 47 -2.01 0.02 47

[Fe ii/H] -2.06 0.06 6 -2.03 0.07 6 -2.01 0.04 8[La ii/Fe i] ... ... 0 ... ... 0 ... ... 0[Mg i/Fe i] 0.53 0.08 3 0.43 0.08 3 0.28 0.07 2[Mn i/Fe i] ... ... 0 ... ... 0 -0.28 0.07 2[Na i/Fe i] -0.08 0.09 2 -0.25 0.12 2 ... ... 0

[Nd ii/Fe i] ... ... 0 ... ... 0 -0.13 0.12 2[Ni i/Fe i] 0.09 0.29 2 -0.01 0.25 1 0.09 0.26 2[O i/Fe i] 0.37 0.15 1 0.32 0.20 1 0.49 0.10 1[Ti i/Fe i] -0.03 0.13 5 -0.13 0.10 4 -0.08 0.06 5

[Ti ii/Fe i] 0.05 0.07 5 0.18 0.09 4 0.16 0.06 5[Y ii/Fe i] -0.10 0.15 4 -0.44 0.23 2 -0.25 0.18 3[Zn i/Fe i] -0.11 0.20 1 0.09 0.20 1 0.11 0.20 1

Cl3-B59 σ Nlines Cl3-B61 σ Nlines Cl3-B82 σ Nlines

[Ba ii/Fe i] 0.27 0.09 3 0.09 0.10 3 0.27 0.09 2[Ca i/Fe i] 0.27 0.17 4 0.21 0.03 6 0.23 0.05 5[Cr i/Fe i] -0.30 0.20 2 -0.28 0.20 2 -0.50 0.15 2

[Eu ii/Fe i] 0.89 0.10 1 0.78 0.15 1 0.97 0.10 1[Fe i/H] -2.35 0.02 44 -2.42 0.03 50 -2.38 0.03 48

[Fe ii/H] -2.30 0.09 5 -2.31 0.09 4 -2.36 0.06 8[La ii/Fe i] <0.52 0.04 5 0.95 0.10 1 0.62 0.06 3[Mg i/Fe i] 0.19 0.09 3 0.37 0.07 3 -0.35 0.08 3[Mn i/Fe i] ... ... 0 0.03 0.07 2 -0.01 0.07 2[Na i/Fe i] 0.05 0.14 2 -0.25 0.12 2 0.48 0.13 2

[Nd ii/Fe i] 0.65 0.07 2 0.44 0.09 2 0.73 0.07 2[Ni i/Fe i] 0.22 0.27 1 -0.04 0.24 1 -0.02 0.23 1[O i/Fe i] 0.43 0.10 1 0.65 0.10 1 0.16 0.10 1[Ti i/Fe i] 0.02 0.08 6 0.04 0.09 4 -0.17 0.07 5

[Ti ii/Fe i] 0.18 0.05 6 0.30 0.08 6 0.05 0.03 5[Y ii/Fe i] -0.23 0.30 2 -0.19 0.19 2 -0.24 0.18 3[Zn i/Fe i] <0.15 0.00 1 <0.22 0.00 1 0.18 0.10 1


Table 5.A3: M15 elemental ratios

M15S1 σ Nlines M15S3 σ Nlines

[Ba ii/Fe i] -0.15 0.03 4 0.29 0.06 4[Ca i/Fe i] 0.32 0.03 7 0.37 0.03 7[Cr i/Fe i] -0.43 0.05 2 -0.33 0.05 2

[Eu ii/Fe i] 0.30 0.20 1 0.65 0.10 1[Fe i/H] -2.36 0.02 55 -2.41 0.02 52

[Fe ii/H] -2.37 0.03 12 -2.35 0.04 12[La ii/Fe i] 0.09 0.10 1 0.36 0.07 3[Mg i/Fe i] 0.33 0.07 3 0.28 0.07 3[Mn i/Fe i] -0.59 0.05 5 -0.38 0.07 6[Na i/Fe i] 0.03 0.14 2 0.68 0.11 2

[Nd ii/Fe i] -0.14 0.12 1 0.30 0.04 2[Ni i/Fe i] 0.07 0.06 3 0.11 0.06 3[O i/Fe i] 0.04 0.10 1 ... ... 0[Ti i/Fe i] -0.08 0.09 5 -0.06 0.07 5

[Ti ii/Fe i] 0.13 0.05 5 0.16 0.07 5[Y ii/Fe i] -0.42 0.02 4 -0.20 0.04 3[Zn i/Fe i] 0.03 0.05 1 0.15 0.09 1

M15S6 σ Nlines M15S7 σ Nlines

[Ba ii/Fe i] 0.39 0.04 4 -0.20 0.06 4[Ca i/Fe i] 0.33 0.03 6 0.41 0.03 7[Cr i/Fe i] -0.39 0.06 2 -0.39 0.04 2

[Eu ii/Fe i] 0.81 0.10 1 0.41 0.10 1[Fe i/H] -2.32 0.03 56 -2.47 0.04 49

[Fe ii/H] -2.35 0.06 13 -2.50 0.04 11[La ii/Fe i] 0.37 0.12 2 0.09 0.18 2[Mg i/Fe i] -0.15 0.07 3 0.43 0.07 3[Mn i/Fe i] -0.41 0.06 5 -0.45 0.20 1[Na i/Fe i] 0.59 0.14 2 0.08 0.12 2

[Nd ii/Fe i] 0.28 0.03 2 -0.23 0.16 1[Ni i/Fe i] 0.06 0.06 3 0.00 0.06 3[O i/Fe i] -0.05 0.10 1 0.15 0.10 1[Ti i/Fe i] -0.09 0.07 6 -0.04 0.11 4

[Ti ii/Fe i] 0.14 0.05 5 0.07 0.07 5[Y ii/Fe i] -0.15 0.04 3 -0.37 0.04 4[Zn i/Fe i] 0.02 0.10 1 0.04 0.08 1

Chapter 6A high resolution spectroscopic studyof Fornax Field Stars

paper in preparation∗

B. Letarte, V. Hill, E. Tolstoy, and DART

ABSTRACT– In this chapter, we present the results of our high reso-lution abundance analysis of 81 individual stars in the central region ofFornax. Using the FLAMES/GIRAFFE spectrograph on the VLT, weobtained high resolution (R ∼20 000) spectra for 81 Red Giant Branchstars in the central 25′ of the Fornax dSph (see chapter 4 and Table 4.3for a description of the observations). Chapters 3 and 4 describe themethods we used to determine abundances, and in this chapter wepresent the sample selection and discuss the abundances that we haveobtained, including α-elements (Mg and Ca, Si, Ti, O), iron-peak (Fe,Ni and Cr) and heavy (Y, Ba, Eu, La and Nd) elements. This is con-sistent with the fact that we randomly selected our sample from theRGB and that the more metal rich stars are centrally concentrated(e.g. Battaglia et al. 2006). We compare our results with Milky Way(MW) studies, and to recent VLT/UVES abundance determinations ofnine individual stars in Fornax globular clusters (chapter 5 and Letarteet al. 2006) and to the Sculptor dSph (Hill et al. in prep). Fornaxstars are found to have unusually low α-elements ratios, as well as Niand Na abundances. The role of metal poor AGB in the creation ofs-process elements is clearly seen by our high [Ba/Y] compared to theMilky Way.

∗ Based on FLAMES observations collected at the European Southern Observatory, proposal number171.B-0588

106 chapter 6: HR spectroscopic study of Fornax Field Stars

Figure 6.1: The spatial distri-bution of the HR spectroscopictargets in the central region ofFornax. The squares are starsfor which we have successfully de-rived abundances and circles forthose that were rejected duringour analysis. The ellipses repre-sent Rc, the core radius and 2 Rc

(Battaglia et al. 2006).

6.1 Sample selection

As first illustrated in Figure 1.3 in chapter 1, our FLAMES HR survey covers thecentral 25′ region, see Figure 6.1. The radial velocities of our targets are presented

in chapter 4, Table 4.4, and the coordinates and photometrically determined Teff of indi-vidual stars are presented in Table 4.A3. Our 107 spectroscopic targets are representedby squares and circles; the squares for stars for which we obtained reliable abundancesand the circles for those that were rejected during our analysis. Stars were rejected forvarious reasons: their Vrad suggested that they are not members of Fornax (1 star); veloc-ity offsets in spectra taken at different times making it impossible to stack the differentexposures (possibly these are binary stars, 8 stars); or the [Fe/H] coming from differentlines were too scattered to perform a conclusive analysis because they are not RGB stars(possibly foreground dwarfs) and/or have a too low signal to noise (17 stars).

Figure 6.2 shows a Colour Magnitude Diagram of a circular region of 10′ radius in thecentre of Fornax, matching the FLAMES field of view, on which we identify our targets,including the rejected ones. As can be seen, we selected our stars to include the entireRGB colour range, going as far to the blue and red side as possible and thus (hopefully)the entire age and metallicity range. We cross-correlated our potential target list withknown carbon stars∗ in Fornax so we could minimise the number of AGB stars thatwe would observe. Still, we were expecting that the bluest and reddest stars might beforeground dwarf or AGB stars but we are confident we didn’t exclude extremely metalpoor stars through our sample selection and subsequent analysis. This is supported byour single metal poor star BL085 at [Fe/H] = -2.58 which was analysed without anyproblems.∗ We are grateful to Serge Demers for providing us with his list

6.2: Results 107

Figure 6.2: A CMD of the FLAMES field of view in the centre of Fornax coming fromour WFI data. Our spectroscopic targets are represented by squares for stars for whichwe were able to determine abundances and circles for those that were rejected during ouranalysis.

6.2 ResultsIn this section, we present the abundance ratios for the 81 stars which survived theselection process described in section 6.1 The stellar parameters we assigned to eachstar, following the method presented in chapter 4 can be found in Table 6.A1. The linelist used, along with the measured EW s for each star (for every element) can be foundin Table 6.A2 and all the abundance ratios used throughout this chapter are listed inTable 6.A3.

6.2.1 Iron abundanceOne of the most straightforward results is an accurate determination of the iron abun-dance, [Fe/H], which in our case is derived from ∼40+ lines per star. We present our[Fe/H] distribution in Figure 6.3 and compare it to the distribution of Battaglia et al.(2006), where they used low resolution (R = 6500) Ca II triplet measurements to de-termine [Fe/H]. Our distribution peaks sharply at [Fe/H] ' -0.8 and is clearly skewedtowards more metal rich stars. There are a few stars outside the main distribution (onthe metal poor side) but the centre of Fornax is clearly dominated by stars with [Fe/H] >-1. However, we still sample two orders of magnitude in [Fe/H], from -2.5 . [Fe/H] . -0.5.

As seen in Figure 1.3 (chapter 1), our FLAMES field included Fornax globular clus-ter 4 but there is no sign that our field sample includes any star from GC 4. This can beseen from the spatial distribution and abundance measurements. The single, really metalpoor star we have in our sample ([Fe/H] ' -2.5) is the only star to overlap with Fornax


Figure 6.3: The [Fe/H] distribution of our Fornax sample (solid line) of stars, with apeak at [Fe/H] ≈ -0.8, richer than the sample of Battaglia et al. (2006) (dashed line).

GCs metallicity range and it is not spatially close to GC 4. Therefore, it seems likelythat it is a metal poor field star, the only representative of Fornax oldest population inour field star sample.

6.2.2 Alpha ElementsThe evolution of chemical abundances in a galaxy is linked to its star formation history,SFH, (e.g. Tinsley 1979; Pagel 1997) and the [α/Fe] ratio is a useful tool to study thisevolution. Alpha elements, which include Calcium (Ca), Magnesium (Mg), Titanium(Ti) Oxygen (O) and Silicon (Si) are predominantly produced by high mass, (> 8 M�)short lifetime Type II supernovae explosions (SNe II). The ratio [α/Fe] is a way of tracingthe relative contribution of SN II and SN Ia products that were available when the starformed. Lower mass SNe II (8-12 M�) result in lower yields than more massive SNe II(Woosley & Weaver 1995). The stars that form after the ISM has been enriched by aSNe II event should have an enhanced [α/Fe] while those that form after the SNe Ia havestarted to contribute significant amounts of Fe to the ISM should have a lower [α/Fe].

The [Fe/H] at which the [α/Fe] ratio starts decreasing (the “knee”) in a galaxy de-pends on several factors: the SFH; the initial mass function (IMF); the time it takes forthe first SN Ia to explode and the time it takes for the mixing of SNe Ia and SNe IIproducts back into the ISM (e.g. Matteucci 2003). The arrival of the first SN Ia shouldbe constant in age (commonly believed to happen after ∼1 Gyr) for all galaxies since it’sthe result of the binary interaction between two evolved stars. The higher in [Fe/H] this“knee” occurs, the more efficient the system was at enriching its gas before the arrival ofthe first SNe Ia. A plateau is reached when (presumably) there is balance between thecontribution of SNe II and SNe Ia.

6.2: Results 109

Figure 6.4: [α/Fe] as a function of [Fe/H]. We computed α as the average of Mg, Caand Ti i abundances. The Milky Way points were taken from the compilation of Vennet al. (2004), with different symbols for thin disk (dark grey squares), thick disk (lightgrey triangles), halo (empty diamonds) stars and our Fornax Field star results (blackcircles). In subsequent plots in this chapter, there will be no distinction between theMW components, for clarity.

In the Milky Way, α-elements are typically overabundant in GCs and metal poor halostars relative to disk stars. To obtain an average measure of α-element abundance, theaverage of Mg, Ca and Ti is taken and is shown in Figure 6.4, where we compare MWhalo, thin and thick disk stars to Fornax Field stars. It is obvious from Figure 6.4 thatour [α/Fe] are significantly underabundant compared to the MW at the same [Fe/H], asign that the evolutionary history of is Fornax is significantly different from the MW.Note that the Milky Way stars in the [Fe/H] range of Fornax are a mix of halo, thin andthick disk stars, a population possibly not as uniform as what we might expect from thecentre of a dwarf galaxy like Fornax. A simple comparison in [Fe/H] of Fornax versus theMW does not take into account the different evolutionary processes in the MW differ-ent components, but is it obvious that Fornax does not overlap with any part of the MW.

Individual [α/Fe] ratios are presented in Figure 6.5 ([Mg/Fe], [Si/Fe] and [Ca/Fe])and Figure 6.6 ([Ti i/Fe] and [Ti ii/Fe]). Along with our Fornax GC and MW points, weintroduce the eight peculiar∗ halo stars observed by Nissen & Schuster (1997) (NS97).These stars were found to display low [α/Fe] and [Ni/Fe] (along with other chemicalpeculiarities) when compared to normal MW halo stars, and this has been suggestedthat they might have been accreted by the MW from a dwarf galaxy.∗ These stars have unusual kinematics and orbital parameters: large maximum distance from the

Galactic centre (Rmax) and large distance from the Galactic plane (zmax).


Figure 6.5: [Mg/Fe], [Si/Fe] and [Ca/Fe] as a function of [Fe/H]. The Fornax field starsare plotted as solid circles, the galactic stars of Venn et al. (2004) as small grey squares,the Fornax globular clusters of chapter 5 as triangles and the eight peculiar halo starsfrom Nissen & Schuster (1997) as empty squares. There is a representative (average)error bar for the Fornax field star abundances in the bottom right corner of each panel.This is the quadratic sum of [element/H] + [Fe/H], (measurement errors) taken fromTable 6.A3. The same axis scale will be used for every [element/Fe] vs [Fe/H] plotspresented in this chapter, for easy comparison.

6.2: Results 111

Figure 6.6: [Ti i/Fe] and ([Ti ii/Fe] as a function of [Fe/H]. Symbols are defined inFigure 6.5.

It can be seen from Figure 6.5 that globally, the metal poor star and the GCs of Fornaxshow a typical overabundance in α-elements but as the [Fe/H] increases, [α/Fe] decreases,reaching a significant underabundance for the highest metallicities. The [Mg/Fe] are typ-ically close to zero, lower than Galactic stars by at least 0.2 to 0.5 dex and the highestvalues just overlap NS97. The [Ca/Fe] ratios are globally lower than other α-elements,especially for the handful of stars close to [Fe/H] = -1.5. At higher metallicity ([Fe/H]& -0.7), our [Ca/Fe] ratios seems to increase, as if the balance in SN II/SN Ia has beenbroken and suddenly there was less contribution from SN II, which seems unlikely, sinceCa is the only element for which we notice such behaviour. The [Si/Fe] ratios are thehighest for the α-elements (if we do not consider Ti ii, which will be discussed later),with a mean value close to zero. The [Si/Fe] in Fornax agree with the eight peculiar halostars of NS97, the only α-element for which this is clearly the case.

In our Teff range, (. 4400 K) atoms of Ti and O can be trapped in molecules, like TiOand CO, artificially lowering our derived abundances. This is not a problem in hotteratmospheres, where the the molecules will always be broken up into their atomic form.We have tested the influence of molecules on atomic abundances by using a line forma-


Figure 6.7: Iron peak elements in Fornax, [Ni/Fe] and [Cr/Fe] as a function of [Fe/H].Symbols are defined in Figure 6.5.

tion code that takes molecules into account∗ to derive our abundances and compare thesewith our standard analysis. We chose stars to cover the Teff range of the Fornax field stars.

We observed a significant difference (∼0.3 dex) in our O/H ratios, confirming thatat this Teff , there is a significant fraction of oxygen locked in the form of CO. The Oabundance is not only affected by molecules, but it also has a problem with the linesused for this analysis. They are two O lines in our wavelength range, the forbidden lineat 6300 Å and another one at 6363 Å. The more reliable line at 6300 Åis not suitable forFornax stars as the typical Vrad means this matches a telluric absorption line, renderingit unusable for most of our stars. The second one is a very weak line, falling in a regionthat is affected by a Ca i auto-ionisation broad feature at 6362Å, and would deserve muchmore attention than could be awarded in the course of this thesis, to provide an accurateO abundance indicator.

∗ TURBOSPECTRUM (Alvarez & Plez 1998), used in the plane-parallel approximation

6.2: Results 113

In contrast to oxygen, we found that Ti abundances are typically not affected bythe amount of Ti locked in the form of TiO. However, it can be seen from Figure 6.6that [Ti i/Fe] is quite different from [Ti ii/Fe], a sign that there are significant non-LTEeffects, affecting more Ti i abundances (since most of the Ti i lines have a low χex) thanTi ii. However, our Ti i abundances are statistically more reliable than Ti ii, since Ti iwas calculated using on average ∼8-9 lines compared to only ∼2-3 lines for Ti ii. This isreflected in the much larger error bars on Ti ii. If we consider only Ti i, it’s behaviour issimilar to a typical α-element like Ca of Mg.

6.2.3 Iron peak elementsAccording to nucleosynthetic predictions, iron peak elements like Iron (Fe), Chromium(Cr), and Nickel (Ni) are believed to be formed predominantly from explosive nucleosyn-thesis, in SN Ia (Iwamoto et al. 1999; Travaglio et al. 2005). In Figure 6.7, we present the[Ni/Fe] and [Cr/Fe] abundance ratios for Fornax stars as a function of [Fe/H]. Both ratiosbehave in more or less the same way, but Ni has much smaller error bars due to the largernumber of lines available; we have ∼15 lines of Ni per star, compared to only 1 line of Cr.[Cr/Fe] seems to be more scattered than [Ni/Fe] but consistent with the larger error bars.

Our metal poor star (BL085, [Fe/H] = -2.58) has a significantly higher [Ni/Fe] thanthe other Fornax field stars, comparable to MW halo stars and Fornax GC stars. At thislow metallicity, we do not detect as many Ni lines as for the more metal rich stars, only4 instead of ∼15 lines, giving it a much larger error bar than the average shown in thebottom corner of the plot as can be seen in Table 6.A3.

For [Cr/Fe], we performed the abundance determination on an Arcturus spectrum,known to have [Cr/Fe] ' 0.0, similar to other MW stars with [Fe/H] ' -0.5, and with thesame line, we obtain [Cr/Fe] = -0.2. The same analysis was done for the Ni lines and weobtained the expected [Ni/Fe] = 0.0 for Arcturus. This is a sign that the Cr line couldhave an erroneous log gf, leading to lower a abundance. We therefore expect that thetrue [Cr/Fe] could be some ∼0.2 higher than what is currently displayed in Figure 6.7,with an average closer to the [Ni/Fe] value.

At [Fe/H] . -1.5, [Ni/Fe] appears to be Galactic halo-like, while near [Fe/H] ' -1.0,it decreases. This [Ni/Fe] underabundance has also been observed in the eight peculiarhalo stars of Nissen & Schuster (1997) (but not for [Cr/Fe]). Similar to the NS97 stars,our [Ni/Fe] underabundance in Fornax is also accompanied by a moderate decrease in[Na/Fe] and [α/Fe] in this metallicity range. These low values of [Ni/Fe] and [Cr/Fe] can-not be easily explained with our current understanding of nucleosynthesis. The [Ni/Fe]ratios should be zero and constant for all [Fe/H] since the two elements are believed tobe predominantly created in the same production site, SN Ia (Travaglio et al. 2005).To witness this different behaviour is an indication that the production factors for eachiron-peak element are not the same and depend on the evolutionary history of the parentpopulation. Maybe the SNe Ia Ni yields are linearly dependant on the original metallicityof the white dwarf progenitor, as suggested by Timmes et al. (2003), or some elements(like Ni but not Fe) are more affected by winds, causing preferential metal loss in theISM. This result is surprising and will need to be investigated further, as Fornax is not


Figure 6.8: [Mg/Fe] plotted against [Na/Fe]. Unlike two of the Fornax GC stars, thereis no sign deep mixing pattern in the field.

the only galaxy in which this behaviour is observed. This under-abundance patterns in[Ni/Fe] and [Cr/Fe] have been observed in the Large Magellanic Cloud (LMC) disk starsby Pompeia et al. (2006).

This under-abundance of Ni and Cr was not observed by Shetrone et al. (2003), withwhich we have three stars in common. We attribute this difference to a combination ofsystematic effects, where the most important one is the use of a different line list. Seesection 4.4.1 and Figure 4.16 for more detail on the systematics in abundance determi-nation.

6.2.4 Deep-mixing patternThe so-called deep mixing pattern is believed to be caused by self pollution in a starthat modifies the upper atmosphere abundances. Proton-capture nucleosynthesis canconvert O, N and Ne to Na, and Mg to Al in the H fusion layer of evolved RGB stars. Adeep-mixing pattern can thus cause a decrease in O associated to an increase in Na anda decrease in Mg with an increase in Al. This has only been observed in stars belongingto globular clusters, in a range of different galaxies, like the Milky Way, the LMC andFornax (see chapter 5).

6.2: Results 115

Figure 6.9: The Na-Ni relationship. Symbols are defined in Figure 6.5.

In the Fornax field stars, we do not find evidence of deep mixing pattern, see Fig-ure 6.8. Unfortunately, our FLAMES observation wavelength coverage (usable) stops atλ ' 6690 Å, just missing two Al lines at 6696.03 Å and 6698.67 Å. It is therefore not pos-sible to say anything about Al. Also, as mentioned in section 6.2.2, we do not yet havea reliable O abundance for our Fornax stars. However, it is clear from Figure 6.8 thatthere is no enhanced Na. All [Na/Fe] values show a significant underabundance, muchlower than MW stars of similar metallicity, going as low as [Na/Fe] = -1.0 in the extremecases. There are only two outliers that have higher [Na/Fe] but these are compatiblewith normal MW stars, showing no hints of deep mixing. There are some low [Mg/Fe]values in Fornax, but without a high [Na/Fe], this is not deep-mixing. In Fornax onlythe GC stars show deep mixing pattern.

6.2.5 The Na-Ni relationshipWhere does Ni come from? As mentioned in section 6.2.3, Ni comes predominantly fromSN Ia, but at earlier times, before the first SN Ia start to enrich the ISM, it is linked tothe Na production in SNe II. As explained in section 5.6 of Clayton (1983), a correlationbetween Na and Ni is a natural result of nucleosynthesis in massive stars. Timmes et al.(1995) suggest that Na is delivered to the ISM when massive stars explode as SNe IIand the amount of Na produced is controlled by the neutron excess, where 23Na is theonly stable neutron-rich isotope produced in significant quantity during C and O burningstage.


During the SN II event, the elements are photodissociated to protons and neutrons,which will recombine to form 56Ni, which β decays to 56Fe, the dominant isotope ofiron. 54Fe and/or 58Ni can also be produced at this stage, depending on the abundanceof the neutron-rich elements (e.g. 23Na). The amount of 54Fe made is small comparedwith the total yield of iron (dominated by the 56Fe production), but this is the mainsource for 58Ni, the stable isotope of nickel. In summary, the Ni production dependson the neutron excess during the photodissociation of the core during the SN II event,and the neutron excess will depend primarily on the amount of 23Na produced earlier.So the Na-Ni relationship is expected when SN II enrichment dominates. The arrivalof SN Ia can break (or flatten) this relationship, as Ni is produced without Na in thestandard model of SNe Ia (Tsujimoto et al. 1995). But the Na yields are still a matter ofdiscussion, because SN Ia involve binary star interactions, and the outcome is dependenton the accreted star. So it is hard to accurately predict the SNe Ia contribution to theNa-Ni correlation.

This has two implications of interest to us: i) In SNe II, the dominant source ofNi is independent of the dominant source of Fe, possibly allowing an underabundance in[Ni/Fe] at low metallicities, before the SN Ia start to enrich the ISM. ii) If the first SNe IIwere not neutron rich, then there will be a lack of Na (therefore also Ni) created. Thenext generation of stars that will be formed could carry this signature, even at highermetallicity. Our low Ni at high metallicity is possibly the consequence of low Na in earliergenerations, caused by neutron-poor SNe II.

For Galactic stars, it has been observed that stars with low [α/Fe] also have low[Na/Fe] and low [Ni/Fe]. Nissen & Schuster (1997) first found this behaviour for eightpeculiar halo stars. They suggested that these stars might have been accreted from anearby dwarf galaxy (their abundance patterns show that they are chemically differentfrom the MW). Other studies, by Fulbright (2002), suggested something similar for starsat large galactocentric distances. Stephens & Boesgaard (2002) suggest that there mightbe a gradient (0.1 dex over 10 kpc) in our galaxy, reflecting the different local conditionswhere the stars form.

If we look at Figure 6.9, our extremely low [Na/Fe] values go beyond the previousgradient (to lower [Na/Fe]) observed by Venn et al. (2004) (their Figure 5). As they havedone, we have restricted the iron abundance of the stars plotted to -1.5 < [Fe/H] < -0.5,where the Na-Ni relation can be seen. This suggests that the chemical evolution of thismetal rich sub-sample of Fornax stars is relatively similar to the MW, at least in how Nais linked to Ni in this [Fe/H] range, with the NS97 points filling the gap between Fornaxand the MW.

6.2.6 Heavy elementsHeavy elements are those with atomic number Z > 30, like Yttrium (Y), Barium (Ba),Europium (Eu), Lanthanum (La) and Neodymium (Nd). They are neutron captureelements built from elements that are exposed to the high neutron flux. Iron peakelements, like 56Fe, are the most efficient seeds to capture neutrons to create heavierelements. There are two main types of neutron capture, the s-process (or slow process)

6.2: Results 117

Figure 6.10: Heavy elements: [Ba/Fe], [Y/Fe] and [La/Fe] as a function of [Fe/H].Symbols are defined in Figure 6.5.


Figure 6.11: Heavy elements in Fornax: [Nd/Fe] and [Eu/Fe] as a function of [Fe/H].Symbols are defined in Figure 6.5. The [Nd/Fe] points are from Burris et al. (2000).

and the r-process (rapid process). When a seed accumulates neutrons and leads to theproduction of a β-unstable nucleus, like 59Fe, the outcome will depend on neutron-capturetimescales. In the s-process, most of the unstable nuclei will have time to undergo β-decay before capturing other neutrons and building up heavier elements (Pagel 1997).The main contribution of the s-process is believed to be from thermal pulses in 2-4 M�AGB (cool giants) stars (Truran 1981). The r-process by contrast usually involves moreextreme conditions, very high temperature and neutron densities, possibly found in lowmass (8-12 M�) SN II explosions.

Tracing the s-process and r-process contribution

We present the abundance ratios of [Ba/Fe], [Y/Fe], [La/Fe], [Nd/Fe] and [Eu/Fe] as afunction of [Fe/H] in Figures 6.10 and 6.11 in decreasing order of s-process contribution.Ba is the element with the larger s-process contribution (for the main∗ s-process) in theSun, with a fraction of 88% s-process (Kappeler et al. 1989). Then, Y has an s-process∗ Contribution from the weak s-process (helium burning in massive star) is 4% for Y, 1% for Ba, La,

Nd and 0% for Eu.

6.2: Results 119

Figure 6.12: [Ba/Y] as a function of [Fe/H]. Symbols are defined in Figure 6.5.

contribution of 85% (Raiteri et al. 1992), The La s-process fraction is 75%, Nd 46%(Kappeler et al. 1989) and only . 5% for Eu, as it is 95% r-process (Burris et al. 2000).Their relative abundance is thus a good discriminant of the dominant neutron captureprocesses in Fornax stars.

According to nucleosynthesis calculations based on hydrodynamical simulations, thes-process does not occur before [Fe/H] & -2.0 and is not significant before [Fe/H] ' -1.0(Travaglio et al. 1999, 2004). This is consistent with what we observe in Fornax, where[Ba/Fe] increases significantly only at [Fe/H] & -1.0. This is different than what weobserve for [Y/Fe], also a s-process-dominated element. From Figure 6.10, we see thatthe [Y/Fe] ratio distribution is flat, scattered around zero in almost the same way asMW stars. This means that the s-process-enrichment does not uniformly contribute tothe creation of different s-process-elements which, in the Sun, have the same s-processcontribution. The behaviour of [La/Fe] is similar to [Ba/Fe], except that the rise in[La/Fe] with increasing [Fe/H] is not as prominent, probably due to the smaller con-tribution from s-process to La. The same reasoning goes for [Nd/Fe] in Figure 6.11:with a smaller s-process contribution comes a smaller increase (barely noticeable) withincreasing [Fe/H]. And finally, [Eu/Fe] shows no increase at all with [Fe/H], confirmingthat Eu is r-process dominated. In summary, from Figures 6.10 and 6.11, we observethat the s-process becomes clearly dominant at [Fe/H] & -1.0.

Dominant role of metal poor AGB

[Ba/Y] behaves quite differently in Fornax than in the MW, as can be seen in Figure 6.12,which clearly shows that Fornax favoured the creation of Ba over Y compared to the MW.Both elements are supposed to be s-process dominated from our knowledge of s-processcontribution in the Sun. Ba (Z=56) belongs to the 2nd peak in the distribution of neutronmagic numbers and Y (Z=39) belongs to the 1st peak. These magic numbers are relatedto low neutron-capture cross-sections which lead to abundance peaks close to Sr (Z=38,next to Y), Ba and Pb (Z=82).


Figure 6.13: [α/Eu] as a function of [Fe/H]. Symbols are defined in Figure 6.5.

The high [Ba/Y] in Fornax could be explained in a scenario where Fornax has a largercontribution from metal poor AGB stars compared to the MW. This would mean thatthe AGB that created the s-process elements were more metal poor in Fornax, favouringthe creation of heavier elements, bypassing first peak elements (Y) in favour of secondpeak elements (Ba). A more metal poor environment has fewer nuclei (seeds) to absorbthe neutrons available from the AGB envelop so they have to accumulate on fewer targets(creating more high-Z elements). In a more metal rich environment, the neutrons willhave more targets on which to distribute, creating more low-Z elements.

6.3 DiscussionIn summary, because it reached low [α/Fe] values at much lower [Fe/H] than the MW,Fornax was most likely less efficient in enriching its gas, possibly due to galactic winds.The lack of Fornax stars (in our sample) with -2.5 < [Fe/H] < -1.0 prevents us frommaking a more conclusive statement, as we cannot trace the decline (the knee) from highto low [α/Fe]. The [α/Fe] deficiency at [Fe/H] & -1.0 could be explained as the resultof an early burst of star formation in the evolutionary history of Fornax followed by adormant period with predominantly SN Ia enrichment before the relatively recent starformation episode that formed most of the stars in our sample (1-3 Gyr ago).

The underabundance of [α/Fe] in Fornax with respect to the MW could be attributedSNe Ia starting to contribute Fe at lower metallicities in Fornax, resulting in a lower ra-tio. Also, perhaps because Fornax is a small galaxy with a low average star formationrate, α-elements were most likely only produced by low-mass SNe II (8-12 M�) whichresult in lower yields than more massive SNe II (Woosley & Weaver 1995). This suggestan “effectively” truncated IMF. This is plausible since intuitively, a small system likeFornax is less likely to form giant molecular clouds (that are believed to be required toobtain massive stars) than a much larger system like the MW.

6.3: Discussion 121

Figure 6.14: [α/Fe] ratio as a function of [Fe/H], including points for the Sculptor dSph(Hill et al. in prep.).

It has been argued that the r-process production occurs in low mass SNe II (Mathewset al. 1992). The high [Eu/Fe] ratios in Fornax (reaching slightly higher than the MW)suggest an important contribution of these low mass SNe II in the enrichment of Fornax.This is compatible with low [α/Fe], for which we also need low mass SNe II. When weplot [Eu/α] ratios (Figure 6.13), it is clear that the sites and relative contribution ofα- and r-process-elements creation differ in Fornax and the (average) MW, where theFornax values are ' 0.7 dex lower than the MW. If both r-process and α−elements werecreated in the same way in every galaxy, the ratio of [Eu/α], should be the same in everysystem. But we know this is not the case, as shown in chapter 5, the stars of Fornaxcluster 3 are r-process enhanced, (Eu-rich) something not common but still seen in manyM 15 stars (Sneden et al. 1997). The reason why all Fornax stars have low [α/Eu] is notthe same for the field stars and the GC. The stars of Cluster 3 are Eu-rich and the fieldstars are α-low.

6.3.1 Comparison of Fornax and SculptorIn Figure 6.14 we compare the average α-element abundance in Fornax with those foundin a similar FLAMES study of the Sculptor dSph (Hill et al., in prep). Sculptor is a faintdSph galaxy (LFnx ' 7×LScl) dominated by old (10-12 Gyr) stars. Using the combineddata set we clearly see a decline in [α/Fe] as SN Ia become more important. It appearsthat the stars observed in Fornax (except for the single metal poor star) are all lyingin the “plateau” where a balance has been achieved between SN Ia and SN II elementproduction. To understand if this is a recent minimum or the end product of the entireevolutionary history of Fornax we need to observe more stars in Fornax in the metallicityrange -2.0 < [Fe/H] < -1.0. This would also allow us to make a more detailed comparisonbetween the metal enrichment history of Scl and Fnx, which are two galaxies with verydifferent star formation histories (cf. Tolstoy et al. 2003).


Figure 6.15: [Ba/Eu] ratio as a function of [Fe/H], including points for the SculptordSph (Hill et al. in prep.).

In Figure 6.15, we compare [Ba/Eu] in Fornax to the same Sculptor results. Using[Ba/Eu] we track the slow increase in the s-process contribution with increasing [Fe/H].As the stellar population becomes more metal rich there is a steady rise in the Baabundance. In Figure 6.15 we also show the Galactic measurements, and it can be seenthat the s-process is a much stronger contribution to chemical evolution of Fornax thanit is to the MW, with a startling divergence at [Fe/H] ' -1.0. This suggests that stellarwinds (e.g., from AGB stars) have played a uniquely important role in the (recent, 2-4 Gyr ago) enrichment history of Fornax. It is also clear that Sculptor with it’s muchshorter star formation history never reached this stage where stellar winds from evolvedstars had a strong effect on the enrichment of the ISM.

6.3.2 Age and [Fe/H]In Figure 6.16 we show a Colour-Magnitude diagram where we have colour-coded thestars depending upon their spectroscopically determined [Fe/H]. From this plot we cansee that the metallicity distribution is strongly peaked, with most stars in a relativelysmall metallicity range which are quite spread out in colour over the RGB. Broadlyspeaking the high metallicity stars (the darker points) are redder than the low metal-licity stars (lighter points). However, it is clear that there are exceptions, and that itcan be risky to use the RGB to determine the metallicity of a complex stellar population.

Using isochrones of the appropriate metallicity we can determine ages for each star(see Battaglia et al. 2006, for details). In Figure 6.17 we plot the ages and metallicitiesfor both the HR (this work) and the LR (Battaglia et al.) [Fe/H] measurements of Fornaxfield stars, including the age-metallicity relation (dashed line) determined by Battagliaet al.. We compare our observations to the age-metallicity relation of Gallart et al.(2005), which is determined from a detailed Colour-Magnitude Diagram analysis. It isnot easy to make strong conclusions on the basis of this comparison, as the spectroscopicsamples consist of only RGB stars (which are thus always more than 1 Gyr old), whereas

6.4: Conclusions 123

Figure 6.16: (left): Colour-Magnitude Diagram of the central (10′ radius) region ofFornax, with a box around the area where we selected our targets. (right): Close-up ofthe CMD where our targets are colour coded (grey-scale) according to their metallicityin the following way. Black for stars of -0.5 > [Fe/H] > -0.7, dark grey for -0.7 > [Fe/H]> -0.9, light grey for -0.9 > [Fe/H] > -1.1 and white for -1.1 > [Fe/H] > -2.6.

the CMD analysis also includes stars as young as ∼200 Myr old. However, it is clearthat, although there is a large spread in [Fe/H] especially at relatively young ages, thespectroscopic measurements suggest that the majority of 1.5-2 Gyr old stars in Fornaxare more metal rich than the chemical evolutionary history from CMD analysis suggeststhey should be. This discrepancy needs further detailed investigation to understand whatis going on. It might be that this highlights a very complex metallicity distribution in theyounger stars in Fornax, possibly as the result of a recent merger with a galaxy havingquite a different (higher) average metallicity. This would likely confuse the determinationof a chemical evolution history from CMD analysis.

6.4 ConclusionsThe young metal poor stars of Fornax have low [α/Fe] ratios associated to a sub-zero[Ni/Fe]. The [α/Fe] dependence on [Fe/H] is different from the Milky Way, showing adifferent efficiency in gas enrichment. Fornax is dominated by s−process at high metal-licity, showing the strong role of (metal poor) AGB in its evolution. The oldest mostmetal poor field star in Fornax is almost indistinguishable from the Galactic halo starsat the same [Fe/H], except in the case of [La/Fe] which may indicate a problem in ourLa measurements rather than anything more fundamental. We expect that the hyperfinestructure correction (which we have not applied yet) for La will not lower it enough tobe compatible with MW stars. There is also a near perfect agreement between the abun-dance ratios observed in individual stars in Galactic globular clusters and in the Fornaxglobular clusters. However, observations of more low metallicity stars will be needed tobetter understand the full chemical evolution history of Fornax dSph.


Figure 6.17: Age and metallicity (Z = 10[Fe/H] × 0.02) of the 81 stars from the HRsample (diamonds), and the LR sample of Battaglia et al. (2006) (empty squares) withthe average value for each age bin (dashed line). The age-metallicity relation (filled graysection) was taken from Gallart et al. (2005),

An interesting next step is to link the stellar abundances in the nearby dwarf galax-ies, especially the old, low metallicity stars to studies carried out at high redshift. Forexample the damped Lyα systems (DLAs) give us similar information about chemicalabundances of the gas in distant systems. These objects are potentially the precursorsof what we observe in the Local Group today.

Further work on Fornax will be to investigate the different regions of this surprisinglycomplex dwarf galaxy in more detail. We would benefit from obtaining more resolutionhigh abundances of individual stars. Although 81 stars is a dramatic improvement on theprevious high resolution study (of 3 field stars) it also leaves several intriguing questionsunanswered. Specifically we would like to study more metal poor field stars in this galaxyin both the central region and the outskirts. It would also be interesting to continue ourdetailed abundance studies of the globular clusters, obtaining high resolution spectra forstars in GC 4 & 5. Another intriguing open question is the nature of the “shell-like”structures found in and around Fornax, and high resolution abundances of individualstars may shed some light on their origins.


6.A:

Largetables

125

Table 6.A1: Stellar parameters of the model used for our sample of 81 stars, along with the minimum EW used for each star.Star Teff log g [Fe/H] vt EWmin Star Teff log g [Fe/H] vt EWmin Star Teff log g [Fe/H] vt EWmin

BL038 3980 0.69 -0.88 2.2 14.24 BL148 4023 0.72 -0.62 2.2 20.62 BL218 3939 0.67 -0.62 2.0 16.17BL045 4122 0.85 -1.09 2.1 12.15 BL149 4100 0.88 -0.91 2.2 14.30 BL221 4056 0.81 -0.86 2.1 13.80BL052 3997 0.72 -1.02 2.4 14.24 BL150 4025 0.80 -0.83 2.2 15.29 BL227 4046 0.84 -0.87 2.0 15.40BL065 4330 0.97 -1.43 1.9 11.33 BL151 4024 0.81 -0.86 2.1 13.37 BL228 3992 0.71 -0.88 2.3 13.31BL076 4065 0.83 -0.86 2.2 13.04 BL155 4060 0.90 -0.74 2.3 17.21 BL229 4014 0.80 -0.71 2.3 15.73BL077 4026 0.80 -0.79 2.2 12.71 BL156 4099 0.89 -1.14 2.2 13.48 BL233 4048 0.83 -0.68 2.2 14.52BL079 4036 0.76 -0.56 2.1 17.98 BL158 4078 0.87 -0.87 2.0 15.79 BL239 4123 0.89 -0.88 2.1 13.80BL081 4062 0.82 -0.62 2.1 14.08 BL160 4027 0.84 -0.87 2.2 14.30 BL242 4063 0.85 -1.04 2.1 12.98BL084 3968 0.72 -0.82 2.1 12.10 BL163 4124 0.88 -0.73 2.3 15.62 BL247 4032 0.85 -0.82 2.3 18.64BL085 4291 0.87 -2.58 2.2 9.18 BL166 4086 0.84 -0.89 2.3 14.24 BL250 3944 0.65 -0.67 2.2 19.14BL091 4162 0.86 -0.96 2.1 11.99 BL168 4113 0.83 -0.88 2.1 15.40 BL253 4003 0.82 -0.66 2.3 14.46BL092 3961 0.74 -0.95 2.0 15.23 BL171 4048 0.87 -0.90 2.4 15.79 BL257 3994 0.78 -0.58 2.3 16.77BL096 4010 0.75 -0.75 2.1 17.98 BL173 3988 0.85 -0.78 2.3 17.77 BL258 4030 0.85 -0.56 2.2 16.88BL097 4060 0.82 -0.92 2.3 12.98 BL180 4114 0.77 -0.90 2.2 13.09 BL260 4009 0.79 -0.85 2.4 14.90BL100 4044 0.84 -0.92 2.1 12.04 BL185 4098 0.69 -0.73 2.1 15.84 BL261 4046 0.82 -0.79 2.0 16.99BL104 4013 0.77 -0.96 2.2 13.37 BL190 3979 0.82 -0.79 2.3 13.37 BL262 4130 0.84 -0.78 2.1 15.68BL113 4187 0.83 -0.75 2.2 13.91 BL195 4261 0.91 -1.00 2.1 11.99 BL266 4212 0.83 -1.44 2.0 12.10BL115 4116 0.79 -1.44 2.0 11.77 BL196 4015 0.76 -1.02 2.4 11.93 BL267 4201 0.80 -0.72 2.1 15.62BL123 3993 0.71 -0.97 2.3 11.66 BL197 3956 0.68 -0.89 2.2 14.19 BL269 3990 0.75 -0.81 2.0 13.86BL125 4080 0.79 -0.73 2.1 13.64 BL203 4037 0.79 -0.83 2.1 15.07 BL278 4072 0.64 -0.72 2.3 15.51BL132 3909 0.65 -0.85 2.1 13.75 BL204 4139 0.97 -1.00 2.4 17.54 BL279 4272 0.95 -1.52 1.6 13.31BL135 4058 0.83 -0.95 2.2 16.28 BL205 4243 0.96 -0.69 2.2 14.13 BL295 3980 0.70 -0.70 2.3 17.38BL138 3939 0.71 -1.01 2.3 14.85 BL208 4159 0.89 -0.66 2.1 13.59 BL300 3990 0.71 -0.92 2.2 15.68BL140 3992 0.75 -0.86 2.1 14.52 BL210 4062 0.81 -0.76 2.2 15.62 BL304 3950 0.70 -0.89 2.3 14.41BL141 4078 0.84 -0.82 2.1 13.64 BL211 4058 0.69 -0.65 2.1 14.13 BL311 4027 0.79 -0.78 2.2 16.94BL146 4076 0.84 -0.92 2.3 12.98 BL213 4032 0.78 -0.86 2.2 13.53 BL315 4139 0.86 -0.81 1.8 16.72BL147 4194 0.94 -1.38 1.8 12.21 BL216 3998 0.75 -0.72 2.2 14.08 BL323 3881 0.66 -0.88 2.3 14.46


Table 6.A2: Complete line list with parameters and associated EW s (in mÅ, measured by DAOSPEC) for all the stars. Part 1/5.

Equivalent witdh, one star per column, BLxxxλ(Å) elem χ gf 038 045 052 065 076 077 079 081 084 085 091 092 096 097 100 1046141.73 Ba ii 0.70 -0.08 265.7 182.3 227.5 174.2 236.8 226.5 290.2 267.4 238.8 141.6 226.5 224.8 248.4 247.7 197.9 238.56496.91 Ba ii 0.60 -0.38 259.7 204.4 278.5 161.6 257.9 244.2 ... 287.1 272.3 134.5 231.4 246.9 260.8 274.8 228.3 252.46122.23 Ca i 1.89 -0.32 243.6 219.3 218.2 176.0 242.1 237.9 277.5 265.2 239.3 116.5 214.9 ... 261.2 ... 213.9 236.96156.03 Ca i 2.52 -2.39 37.9 ... 23.5 ... 28.3 19.7 52.0 27.6 17.0 ... ... 24.2 36.5 25.5 ... 22.96161.30 Ca i 2.52 -1.27 144.3 97.1 106.5 44.0 112.7 121.4 167.2 142.5 130.2 17.2 119.5 117.6 141.9 123.9 103.5 131.96162.17 Ca i 1.90 -0.32 271.8 248.2 249.9 200.2 266.1 252.9 295.7 293.0 260.4 135.6 243.2 262.4 284.8 267.7 235.3 260.86166.44 Ca i 2.52 -1.14 141.2 97.5 107.2 51.6 126.1 114.7 134.7 132.2 122.5 ... 105.6 121.1 123.3 119.1 99.9 122.56169.04 Ca i 2.52 -0.80 125.4 121.0 120.8 91.9 146.0 143.9 128.0 151.6 144.6 22.7 134.4 130.2 ... 146.3 115.4 139.06169.56 Ca i 2.52 -0.48 149.8 131.0 135.8 112.8 145.6 164.6 160.8 174.6 158.9 35.6 153.7 156.5 ... 161.8 138.9 166.06439.08 Ca i 2.52 0.39 234.7 199.4 229.8 167.9 228.2 227.2 258.3 237.6 219.9 105.8 205.8 215.8 ... 228.6 201.2 217.96455.60 Ca i 2.52 -1.29 146.3 102.4 31.4 117.2 76.1 ... 40.6 148.5 134.7 ... 123.3 85.0 126.2 154.5 48.1 127.36471.67 Ca i 2.52 -0.76 ... ... ... ... ... ... ... ... ... 20.7 ... ... ... ... ... ...6493.79 Ca i 2.52 -0.32 217.7 166.2 194.6 143.0 177.2 189.8 239.5 195.4 186.5 78.3 169.0 177.1 181.6 174.1 175.1 179.16499.65 Ca i 2.52 -0.82 137.9 101.5 130.4 59.6 142.0 117.8 165.6 136.4 137.1 25.3 118.4 133.8 153.3 118.0 124.2 127.46508.84 Ca i 2.52 -2.41 38.6 ... 24.6 ... 33.7 22.4 67.9 43.1 26.6 ... ... 37.0 ... 37.3 23.1 38.86330.09 Cr i 0.94 -2.92 152.6 108.7 130.8 35.1 133.9 128.0 164.9 139.9 134.6 ... 115.6 131.3 137.0 122.6 120.7 106.46645.13 Eu ii 1.37 0.20 85.0 27.3 81.5 23.3 50.9 50.7 90.1 59.0 65.4 ... 57.9 62.2 79.0 60.7 ... 76.05369.96 Fe i 4.37 0.54 146.2 154.2 125.0 ... 175.2 177.0 ... 180.7 173.9 41.0 159.5 147.4 162.6 145.3 163.9 143.05383.37 Fe i 4.31 0.50 165.1 143.0 158.5 126.9 162.0 157.5 154.3 163.4 153.7 68.1 164.4 141.6 166.2 167.7 150.0 158.55386.34 Fe i 4.16 -1.74 37.7 62.7 30.2 31.0 ... 54.6 ... 33.9 66.2 ... 32.3 45.6 ... 45.1 39.1 ...5393.17 Fe i 3.24 -0.92 191.0 191.9 173.7 149.5 196.3 202.0 210.3 208.5 206.8 77.1 205.9 175.9 189.2 206.5 176.0 195.25395.22 Fe i 4.45 -1.73 ... 22.3 ... ... ... ... ... 41.1 ... ... 21.4 ... ... 23.4 16.7 ...5405.79 Fe i 0.99 -1.85 ... ... ... 268.7 ... ... ... ... ... 205.0 ... ... ... ... ... ...5415.19 Fe i 4.39 0.51 115.7 155.6 119.1 137.7 158.7 166.2 148.2 165.9 161.9 71.0 164.8 142.3 149.4 163.3 154.7 145.35417.04 Fe i 4.42 -1.42 31.2 27.9 27.6 ... 33.1 40.6 31.0 63.1 37.8 24.9 49.4 33.9 28.8 37.1 32.1 46.65434.53 Fe i 1.01 -2.12 ... ... ... 261.2 ... ... ... ... ... 206.0 ... ... ... ... ... ...5436.30 Fe i 4.39 -1.35 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...5464.29 Fe i 4.14 -1.62 74.9 41.2 52.9 ... 51.1 57.9 86.4 72.4 73.4 ... 48.5 54.0 55.0 59.4 57.9 56.75470.09 Fe i 4.45 -1.60 43.2 21.8 26.7 ... 29.5 ... 34.9 41.2 ... ... 18.7 30.8 27.0 25.6 20.8 22.25501.48 Fe i 0.96 -3.05 258.9 236.3 239.1 193.2 247.9 269.1 265.1 276.3 273.9 159.8 241.3 254.3 253.3 267.3 243.0 254.75506.79 Fe i 0.99 -2.79 ... 273.2 294.2 203.6 ... ... ... ... ... 158.7 290.6 ... ... ... 272.1 ...5539.29 Fe i 3.64 -2.59 47.1 41.0 27.4 31.9 57.1 67.1 62.0 82.3 65.6 ... 56.6 42.3 53.9 63.5 54.8 54.85586.77 Fe i 3.37 -0.10 194.0 203.2 141.5 175.9 211.2 232.9 175.6 238.8 226.1 108.8 215.9 207.6 214.7 228.6 221.2 229.26120.26 Fe i 0.91 -5.94 70.3 67.7 92.1 33.9 87.0 85.8 87.4 106.2 83.2 ... 57.9 92.5 94.9 96.2 67.6 91.06136.62 Fe i 2.45 -1.50 178.0 ... 261.1 209.1 ... ... ... ... ... 118.1 ... ... 268.1 ... ... ...6137.00 Fe i 2.20 -2.95 ... ... 181.4 ... ... ... ... ... ... ... ... ... 165.2 ... ... ...6151.62 Fe i 2.18 -3.37 132.7 122.7 137.3 93.6 125.6 148.1 128.0 150.7 126.3 21.4 136.7 133.4 154.4 130.7 126.1 133.96157.75 Fe i 4.07 -1.26 125.5 94.5 102.3 51.0 110.7 107.7 137.9 122.8 112.6 ... 94.7 103.1 120.9 113.0 95.4 115.76159.38 Fe i 4.61 -1.97 47.8 ... ... ... ... ... ... 28.9 ... ... ... ... 27.8 ... ... 25.26165.36 Fe i 4.14 -1.47 64.4 56.4 66.5 37.2 65.2 74.7 88.5 63.9 60.9 ... 66.1 62.7 73.7 68.4 68.0 50.66173.34 Fe i 2.22 -2.85 169.3 150.5 169.6 102.5 160.9 166.9 176.6 176.2 169.7 33.9 148.4 164.9 167.7 164.6 155.9 161.56180.20 Fe i 2.73 -2.78 49.8 119.4 93.6 79.4 119.5 131.8 112.9 ... 134.2 34.7 130.9 116.4 144.3 123.9 112.0 123.56187.99 Fe i 3.94 -1.58 81.4 65.1 72.7 32.1 76.7 81.8 90.9 96.2 78.7 ... 77.9 71.1 76.7 87.8 67.2 79.26200.31 Fe i 2.61 -2.44 138.8 138.5 151.4 114.0 151.2 151.5 147.9 164.4 150.6 39.3 144.0 145.4 157.9 152.0 140.0 158.76213.43 Fe i 2.22 -2.66 188.6 170.0 195.5 143.4 183.2 203.7 207.2 200.8 200.5 72.5 165.1 178.5 185.4 182.4 175.7 180.66219.29 Fe i 2.20 -2.44 199.9 185.7 209.7 135.0 187.8 199.5 203.6 210.5 206.1 68.3 190.2 195.3 205.4 191.4 179.8 190.36226.74 Fe i 3.88 -2.20 60.2 41.4 62.4 21.0 59.3 57.1 74.1 48.4 51.9 ... 58.8 70.3 72.2 55.6 57.4 61.76252.57 Fe i 2.40 -1.76 217.9 193.7 200.3 154.7 207.4 218.5 203.1 214.9 203.2 115.3 187.0 198.0 216.1 205.6 190.3 205.36265.13 Fe i 2.18 -2.55 178.1 167.1 180.8 113.5 179.0 194.9 190.2 186.1 183.0 66.5 170.4 190.9 192.1 196.7 181.2 171.26271.28 Fe i 3.32 -2.96 49.7 46.8 59.8 21.5 59.2 63.2 65.5 72.3 56.8 ... 53.4 49.4 67.4 65.0 44.8 49.26297.80 Fe i 2.22 -2.74 ... ... ... ... ... ... ... ... ... 42.0 ... ... ... ... ... ...6301.50 Fe i 3.65 -0.72 ... 144.9 ... ... ... ... ... ... ... 40.0 ... ... ... ... ... ...6307.85 Fe i 3.64 -3.27 ... ... ... ... ... ... 33.6 26.1 ... ... ... ... ... ... 15.8 ...6322.69 Fe i 2.59 -2.43 163.0 140.4 158.5 116.1 160.4 174.6 169.0 170.1 155.4 36.1 146.2 146.2 150.4 159.7 153.5 165.06330.85 Fe i 4.73 -1.22 37.8 26.3 35.3 17.3 42.6 38.5 41.3 50.7 31.4 ... 34.7 44.3 34.0 35.6 36.5 17.16335.33 Fe i 2.20 -2.23 214.5 198.2 250.2 151.0 224.3 209.1 234.5 220.8 206.2 83.2 197.4 191.7 207.9 213.3 201.1 212.56336.82 Fe i 3.69 -1.05 144.3 144.9 145.5 108.8 140.5 139.9 157.7 144.0 142.3 33.7 150.9 129.9 138.0 134.0 130.4 128.66344.15 Fe i 2.43 -2.92 169.6 165.4 165.2 100.0 170.3 181.3 196.6 189.3 168.3 20.7 149.4 ... 180.6 170.0 143.6 165.76355.04 Fe i 2.84 -2.29 164.2 143.2 149.7 86.5 160.5 163.5 176.3 148.2 153.6 35.2 150.7 159.1 ... 162.0 150.4 160.16380.75 Fe i 4.19 -1.50 79.5 49.9 74.2 30.1 76.5 79.5 88.1 87.6 73.3 ... 72.3 78.5 66.7 73.8 63.3 88.96392.54 Fe i 2.28 -3.95 80.1 84.8 82.9 46.2 85.7 89.3 78.2 106.7 90.0 17.3 78.8 94.2 95.9 87.9 87.5 94.16393.61 Fe i 2.43 -1.63 236.5 228.0 221.5 194.6 241.1 248.4 249.6 260.1 236.6 120.8 240.6 204.9 240.4 251.8 221.3 251.56408.03 Fe i 3.69 -1.00 149.0 140.4 142.1 101.0 149.6 153.4 131.4 167.8 140.8 32.2 141.5 118.5 143.8 142.3 123.3 142.76419.96 Fe i 4.73 -0.24 112.8 91.1 115.3 61.7 99.9 115.5 132.5 118.3 105.1 ... 96.8 76.4 115.4 110.0 109.5 97.16421.36 Fe i 2.28 -2.01 228.9 217.8 233.8 163.9 237.0 234.2 255.5 234.9 223.8 92.9 213.5 218.3 226.9 222.6 214.7 223.86430.86 Fe i 2.18 -1.95 256.5 215.3 282.6 174.9 256.8 269.1 ... 272.7 244.7 111.1 222.4 241.4 276.1 263.7 235.9 242.26481.87 Fe i 2.27 -2.98 183.4 150.6 183.1 119.2 175.7 176.4 210.5 187.6 169.5 39.8 154.5 167.9 169.3 180.6 173.1 182.96498.94 Fe i 0.96 -4.69 205.1 151.9 215.6 110.2 178.6 198.4 243.4 203.5 191.3 28.5 163.7 186.0 208.9 185.0 176.6 205.96533.93 Fe i 4.55 -1.46 43.4 34.9 ... ... 52.3 51.2 70.7 63.9 39.5 ... 40.0 55.9 50.8 32.2 40.7 37.96556.81 Fe i 4.79 -1.72 26.7 20.3 ... ... ... ... 51.0 24.6 17.7 ... ... 28.9 38.5 19.3 21.4 ...6569.22 Fe i 4.73 -0.42 110.3 84.1 122.0 59.8 115.9 114.9 159.8 124.1 106.0 ... 100.0 104.6 110.5 102.6 111.7 100.36574.23 Fe i 0.99 -5.02 191.0 151.3 ... 101.1 177.6 180.7 169.8 191.4 180.5 32.7 150.5 162.1 151.7 186.5 158.2 166.96593.88 Fe i 2.43 -2.39 194.6 167.8 100.8 156.8 179.5 192.2 151.5 220.5 217.3 76.2 179.1 156.4 160.7 200.6 173.5 177.56597.56 Fe i 4.79 -1.07 45.3 44.8 53.7 ... 66.1 66.7 75.1 68.7 64.2 ... 60.3 43.5 56.1 54.4 58.7 48.36608.03 Fe i 2.28 -3.94 92.1 73.6 143.8 35.5 99.6 93.2 154.2 104.1 96.2 29.8 79.7 75.8 119.0 92.3 98.9 94.66609.12 Fe i 2.56 -2.66 158.0 149.2 206.9 107.1 167.1 171.1 208.3 179.3 170.5 36.6 156.9 149.4 184.9 159.7 174.4 151.76627.54 Fe i 4.54 -1.68 16.9 22.7 83.1 ... 29.7 32.9 75.1 43.1 ... ... 25.0 37.4 55.4 25.2 27.6 28.66633.76 Fe i 4.56 -0.82 107.4 86.2 149.8 ... 112.2 94.4 155.8 115.3 85.9 ... 102.2 110.4 ... 107.5 95.4 88.36646.93 Fe i 2.60 -3.99 63.5 42.7 79.8 27.3 66.0 59.1 92.6 62.5 58.8 ... 49.3 67.5 89.5 52.7 59.8 49.26653.85 Fe i 4.15 -2.52 26.6 23.3 69.6 15.0 24.7 24.0 62.2 34.7 36.1 ... 26.0 ... 37.3 28.9 16.5 24.55414.08 Fe ii 3.22 -3.61 ... 23.3 ... 42.9 33.4 20.9 ... ... ... 16.4 34.1 ... 34.4 34.1 20.7 ...5425.25 Fe ii 3.20 -3.36 56.9 50.1 45.7 64.2 44.8 36.0 35.5 38.7 45.3 29.5 53.4 43.7 39.3 37.8 57.1 54.06149.25 Fe ii 3.89 -2.72 36.3 39.3 ... 23.8 ... ... ... 25.0 22.6 ... 21.2 25.9 34.5 36.8 36.1 24.06432.68 Fe ii 2.89 -3.71 37.7 57.7 72.7 41.8 68.6 64.9 75.8 51.9 48.7 16.7 65.8 51.7 ... 55.3 45.5 54.46456.39 Fe ii 3.90 -2.08 56.4 65.0 ... 99.8 83.9 93.6 ... 92.8 64.2 37.7 113.3 ... 31.4 129.0 39.0 105.76320.43 La ii 0.17 -1.56 90.7 43.2 ... ... ... 41.7 101.7 88.1 68.3 18.4 57.3 64.6 80.8 58.0 44.3 76.66390.46 La ii 0.32 -1.40 71.4 ... 57.1 ... 57.3 50.9 93.8 76.1 78.6 ... 57.7 65.1 ... 66.8 49.1 ...5528.41 Mg i 4.35 -0.36 221.1 195.6 212.0 165.9 207.0 200.9 253.7 208.8 217.3 112.8 203.8 210.0 220.0 234.8 198.8 219.0



λ(Å) elem χ gf 038 045 052 065 076 077 079 081 084 085 091 092 096 097 100 1046318.72 Mg i 5.11 -1.97 43.6 ... ... ... 34.9 24.6 ... ... ... ... 31.1 38.9 52.2 ... 17.7 34.06319.24 Mg i 5.11 -2.21 ... 27.0 26.9 ... ... 23.3 36.2 24.9 ... ... ... 30.6 42.1 ... 31.8 ...6319.49 Mg i 5.11 -2.43 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...5420.36 Mn i 2.14 -1.46 217.4 160.5 166.5 67.8 188.6 198.1 204.4 203.5 202.4 ... 176.3 162.5 192.5 209.3 174.0 187.55432.55 Mn i 0.00 -3.80 287.6 247.7 275.0 120.2 280.9 272.6 ... 291.3 252.2 ... 215.8 251.6 289.0 271.4 249.7 264.55516.77 Mn i 2.18 -1.85 169.4 99.3 141.5 36.1 145.4 145.0 179.4 171.6 137.9 ... 117.2 132.7 141.2 141.7 111.8 138.66154.23 Na i 2.10 -1.56 38.8 ... ... ... 23.5 ... 35.7 36.3 ... ... 25.6 15.5 23.4 19.8 ... ...6160.75 Na i 2.10 -1.26 58.4 ... 23.3 ... 38.2 38.6 65.5 54.5 ... ... 38.5 21.1 47.6 39.7 ... 28.55416.38 Nd ii 0.86 -0.98 17.2 ... ... 17.1 27.5 20.6 18.5 34.1 29.3 ... 32.0 25.3 27.5 28.5 16.3 40.05431.54 Nd ii 1.12 -0.47 65.0 ... 43.7 ... 47.3 42.4 73.6 62.0 58.0 35.3 42.5 45.0 72.1 48.3 39.4 52.35485.71 Nd ii 1.26 -0.12 40.7 19.3 25.6 ... 27.4 25.2 48.2 34.7 30.1 ... 45.2 39.3 37.3 39.4 ... 35.95578.73 Ni i 1.68 -2.67 135.0 119.9 116.0 86.8 117.7 130.9 142.9 140.0 138.5 29.0 133.0 123.0 125.5 141.0 115.4 151.55587.87 Ni i 1.93 -2.37 147.1 119.4 121.6 83.3 147.5 159.3 150.8 183.1 157.3 30.6 158.0 129.2 146.5 176.0 135.9 158.75589.37 Ni i 3.90 -1.15 25.0 19.1 26.3 ... 32.4 24.6 29.2 34.7 29.1 ... 18.7 31.3 26.7 18.8 23.9 27.75593.75 Ni i 3.90 -0.79 41.1 31.2 33.4 21.8 34.5 47.1 43.5 57.6 31.9 ... 35.1 40.6 34.2 30.7 24.5 31.16128.97 Ni i 1.68 -3.39 94.8 61.2 93.5 48.1 87.6 76.9 99.4 90.0 82.1 ... 74.3 72.9 105.7 88.8 75.0 82.86130.14 Ni i 4.27 -0.98 17.1 ... ... ... ... ... ... ... ... ... ... 17.9 ... ... 17.1 ...6177.25 Ni i 1.83 -3.60 72.6 40.2 51.7 ... 55.4 57.7 76.8 73.3 65.4 ... 69.5 ... 52.4 72.2 45.0 61.66186.72 Ni i 4.11 -0.90 40.7 28.3 22.8 ... ... 41.8 32.9 43.6 31.0 ... ... ... ... 30.2 26.9 ...6204.61 Ni i 4.09 -1.15 ... ... ... 18.7 20.8 22.3 24.5 22.4 ... ... 28.9 ... ... 23.0 22.4 ...6223.99 Ni i 4.10 -0.97 24.2 ... ... 24.5 ... 30.3 37.9 24.2 ... ... ... 43.0 37.4 30.9 29.2 ...6230.10 Ni i 4.11 -1.20 45.7 23.6 24.5 ... 26.5 ... 42.7 ... 26.7 ... 25.4 ... ... 32.9 ... ...6322.17 Ni i 4.15 -1.21 21.0 ... ... ... ... 27.0 ... 20.4 ... ... ... ... 24.8 19.2 ... ...6327.60 Ni i 1.68 -3.09 110.9 94.9 100.3 70.4 109.7 106.5 127.7 116.7 107.2 ... 102.2 119.9 102.0 117.2 106.6 113.66378.26 Ni i 4.15 -0.82 29.4 ... ... ... 25.9 30.9 ... 34.2 32.7 ... 24.2 19.5 34.5 ... 25.8 43.76384.67 Ni i 4.15 -1.00 22.9 ... ... ... 33.3 ... 50.2 35.4 30.4 ... ... ... 20.9 27.6 ... 36.26482.80 Ni i 1.94 -2.85 110.5 96.8 141.8 48.7 113.0 108.0 147.1 123.6 107.3 16.2 92.5 104.6 115.4 109.2 106.8 107.76586.32 Ni i 1.95 -2.79 137.0 135.4 44.1 123.9 66.9 89.8 ... 134.5 144.6 45.8 128.7 57.0 51.5 148.1 66.5 94.66598.61 Ni i 4.24 -0.93 ... ... ... ... ... ... ... ... ... 17.7 ... ... ... 34.9 24.2 ...6635.14 Ni i 4.42 -0.75 ... 15.1 ... ... ... ... 46.9 ... ... ... 18.6 24.9 ... ... ... 18.26300.31 O i 0.00 -9.75 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6363.79 O i 0.02 -10.25 53.3 26.8 39.1 28.4 37.2 45.8 53.3 45.0 34.4 ... 24.6 33.5 57.3 39.5 30.7 46.35526.82 Sc ii 1.77 0.03 118.4 113.7 124.9 104.2 109.3 107.1 116.3 114.5 119.4 74.4 121.6 110.8 114.1 120.8 107.5 115.46245.62 Sc ii 1.51 -0.97 90.2 58.8 67.4 56.9 65.4 65.0 77.1 58.3 81.7 24.0 72.5 75.4 73.2 68.5 61.8 66.96309.90 Sc ii 1.50 -1.52 ... 33.4 ... 26.3 ... ... 47.8 ... 38.3 ... ... 27.7 42.5 ... 41.7 ...6604.60 Sc ii 1.36 -1.31 87.1 84.3 109.3 61.2 98.5 80.1 125.6 83.5 105.3 38.9 76.4 90.6 89.4 108.5 82.1 80.46125.03 Si i 5.62 -1.57 25.0 24.7 ... 18.3 ... ... ... ... ... ... 22.4 ... ... ... ... 34.26142.48 Si i 5.62 -1.51 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6145.02 Si i 5.61 -1.37 45.1 ... ... ... 16.0 ... ... ... ... ... ... 22.0 ... ... ... 15.26155.14 Si i 5.62 -0.80 42.5 38.5 ... 32.4 35.6 39.2 57.8 31.5 27.3 17.0 35.0 25.7 48.1 38.3 23.9 40.16237.33 Si i 5.61 -1.02 26.6 ... ... ... ... ... 34.7 ... ... ... 21.3 25.5 26.1 24.6 ... ...6243.82 Si i 5.61 -1.27 ... ... 23.2 ... 22.8 24.1 ... ... 28.2 ... ... 24.1 ... 33.2 ... ...5490.16 Ti i 1.46 -0.93 121.8 83.7 104.4 34.0 109.1 102.3 142.0 125.3 112.4 ... 84.4 103.8 109.5 98.8 84.1 116.15503.90 Ti i 2.58 -0.19 ... 41.0 ... 19.8 47.0 47.1 83.9 83.4 55.3 ... 43.9 74.4 56.7 71.7 33.9 58.46126.22 Ti i 1.07 -1.42 149.2 110.7 119.7 48.9 131.6 130.4 167.1 150.6 142.1 20.0 113.5 107.1 158.2 135.5 113.3 140.16220.50 Ti i 2.68 -0.14 78.1 45.9 ... ... ... 48.8 79.3 78.6 ... ... ... ... ... 58.6 ... ...6258.10 Ti i 1.44 -0.35 208.0 139.2 ... 77.3 176.3 159.6 231.6 191.6 195.5 ... 143.9 192.0 ... 182.0 ... 194.36303.77 Ti i 1.44 -1.57 ... 61.1 66.2 26.0 74.9 72.7 111.6 101.7 84.2 ... 50.4 88.8 102.0 80.1 55.4 83.26312.24 Ti i 1.46 -1.55 ... ... 66.8 ... 63.2 66.9 115.9 97.8 79.9 ... 51.4 88.8 ... 80.6 54.5 88.26336.10 Ti i 1.44 -1.74 94.4 48.9 63.7 16.9 74.7 65.8 111.4 84.7 73.2 ... 41.0 72.2 90.8 70.4 47.0 71.66508.12 Ti i 1.43 -2.05 64.6 27.1 ... ... 62.0 40.6 105.0 61.9 50.3 ... 28.1 58.9 59.3 37.8 30.2 58.86556.08 Ti i 1.46 -1.07 ... 104.7 ... ... ... ... ... ... ... ... ... ... 164.6 ... ... ...6599.13 Ti i 0.90 -2.09 151.6 81.6 128.0 35.6 141.1 128.7 179.9 146.9 149.8 ... 87.8 109.4 141.6 130.8 105.4 120.16666.53 Ti i 1.46 -1.62 25.9 ... 15.1 ... ... 18.7 28.2 25.8 ... ... 18.0 ... ... ... ... ...5418.77 Ti ii 1.58 -2.11 111.6 107.9 73.6 89.0 94.6 98.7 97.3 91.1 103.5 69.1 96.9 94.9 108.1 108.4 104.0 103.06219.94 Ti ii 2.06 -2.82 27.2 15.0 ... ... ... ... ... ... ... ... ... ... ... ... ... ...6559.58 Ti ii 2.05 -2.02 ... 48.3 ... 45.2 ... ... ... ... ... 16.4 ... ... 81.2 ... ... ...6606.95 Ti ii 2.06 -2.79 42.2 30.2 ... 25.2 42.1 31.9 70.3 36.3 47.5 20.9 44.5 36.2 46.3 43.1 44.9 35.86680.13 Ti ii 3.09 -1.86 ... ... 23.9 18.9 22.7 ... ... 19.0 ... 16.6 32.6 27.0 ... 16.8 ... 16.36119.53 V i 1.06 -0.32 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 71.86128.33 V i 1.05 -2.30 ... ... ... ... ... ... 30.3 ... ... ... ... ... ... ... ... ...6135.37 V i 1.05 -0.75 126.5 48.6 95.7 18.2 87.3 87.7 122.7 107.7 98.3 ... 73.6 101.1 100.9 87.5 62.6 ...6150.15 V i 0.30 -1.79 156.7 98.6 118.7 18.7 125.6 129.4 186.0 154.2 146.0 17.3 99.4 124.1 156.0 135.7 91.4 150.56199.19 V i 0.29 -1.29 185.0 133.3 152.9 40.1 163.3 179.1 210.6 186.5 171.9 ... 130.0 180.5 192.4 162.0 131.0 179.26216.36 V i 0.28 -0.81 217.5 141.6 169.3 41.8 176.7 173.2 211.0 184.0 184.0 21.2 136.6 168.0 198.2 165.7 142.7 185.66224.51 V i 0.29 -2.01 127.1 77.8 98.9 24.1 115.9 125.0 138.4 136.0 124.2 ... 81.0 123.9 142.2 115.3 85.3 133.16233.20 V i 0.28 -2.07 102.1 64.1 92.4 ... 95.4 89.7 128.6 96.4 90.2 ... 61.2 106.5 108.4 85.4 53.3 97.86243.11 V i 0.30 -0.98 ... ... ... 77.3 ... ... ... ... ... ... ... ... 298.9 ... ... ...6251.82 V i 0.29 -1.30 163.8 100.2 140.9 36.0 137.2 137.9 174.9 157.7 158.0 ... 115.9 137.2 158.8 139.1 109.5 149.06274.66 V i 0.27 -1.67 ... ... 91.9 ... 115.1 116.4 155.2 141.2 ... ... 92.2 131.1 ... 94.0 90.8 141.86357.29 V i 1.85 -0.91 38.3 ... ... ... ... 17.8 43.7 25.7 ... 15.5 ... 23.8 50.3 ... 15.4 21.26452.32 V i 1.19 -1.21 89.2 35.8 32.1 48.4 65.7 72.3 79.2 122.1 85.8 ... 48.5 68.1 59.3 73.6 35.8 74.06504.19 V i 1.18 -1.23 60.1 16.5 53.9 ... 41.6 33.3 104.7 62.9 38.6 ... 32.9 64.6 59.4 49.1 36.6 59.16531.41 V i 1.22 -0.84 74.8 ... ... ... ... ... ... ... ... ... ... ... 70.8 ... ... ...5402.78 Y ii 1.84 -0.51 52.0 ... 30.7 ... ... ... 67.2 35.5 26.3 ... 25.0 52.1 31.9 39.3 20.5 43.36362.35 Zn i 5.80 0.14 ... 17.1 ... ... ... ... ... 36.6 ... ... ... ... 92.9 25.2 ... ...6127.48 Zr i 0.15 -1.06 101.9 33.8 57.2 ... 69.5 67.8 112.9 99.4 79.8 18.2 32.0 95.0 91.7 77.6 43.3 84.56140.46 Zr i 0.52 -1.41 36.1 ... ... ... ... 15.7 30.1 22.6 ... ... ... ... 21.5 ... 18.5 ...6143.18 Zr i 0.07 -1.10 120.7 40.0 62.8 ... 74.5 65.4 129.8 101.7 70.6 ... 40.7 70.9 87.5 73.9 ... 96.06192.95 Zr i 0.54 -2.07 31.8 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...



Equivalent witdh, one star per column, BLxxxλ(Å) elem χ gf 113 115 123 125 132 135 138 140 141 146 147 148 149 150 151 1556141.73 Ba ii 0.70 -0.08 276.9 182.4 241.5 256.1 218.8 240.3 255.0 241.1 209.2 240.3 236.5 ... 236.0 230.1 243.6 265.26496.91 Ba ii 0.60 -0.38 ... 171.9 235.9 277.9 248.3 255.0 227.5 257.0 227.5 266.2 244.7 227.2 257.8 236.3 253.6 293.66122.23 Ca i 1.89 -0.32 249.6 192.4 242.8 254.1 250.0 244.7 255.0 234.6 234.5 230.0 162.3 293.5 228.1 256.4 240.9 ...6156.03 Ca i 2.52 -2.39 43.4 ... ... 40.5 28.5 ... 24.2 34.5 ... ... ... 70.9 18.7 36.6 22.4 27.66161.30 Ca i 2.52 -1.27 156.4 52.9 131.7 145.8 126.1 113.7 148.3 131.4 119.2 116.1 66.5 170.4 111.5 125.8 119.4 155.36162.17 Ca i 1.90 -0.32 264.2 204.5 262.9 284.9 261.6 266.1 279.5 254.3 247.1 265.4 194.6 ... 244.5 268.6 265.7 277.26166.44 Ca i 2.52 -1.14 127.7 60.6 111.8 124.8 121.9 119.2 145.6 133.1 105.3 117.6 52.9 145.8 112.6 116.7 124.8 138.26169.04 Ca i 2.52 -0.80 145.9 87.2 141.7 139.1 135.1 129.7 153.1 148.0 149.2 139.0 59.7 143.1 116.5 155.5 148.3 147.96169.56 Ca i 2.52 -0.48 165.6 111.5 157.5 178.6 155.6 175.7 146.1 154.7 156.3 167.7 90.4 174.5 144.9 167.7 158.0 187.96439.08 Ca i 2.52 0.39 232.7 168.5 225.1 235.0 226.6 235.2 238.2 214.9 219.1 227.5 181.8 253.0 230.3 240.7 220.4 251.16455.60 Ca i 2.52 -1.29 155.8 85.6 106.2 92.6 172.9 192.2 134.1 135.9 159.0 155.3 ... 178.6 ... 194.4 177.9 149.06471.67 Ca i 2.52 -0.76 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6493.79 Ca i 2.52 -0.32 195.3 139.0 193.3 213.1 187.8 174.8 196.5 176.7 192.2 186.3 134.8 205.0 188.9 199.2 183.4 193.16499.65 Ca i 2.52 -0.82 129.7 80.5 132.1 147.0 129.4 122.0 130.1 136.8 128.4 118.6 80.8 139.6 122.0 120.3 134.2 133.46508.84 Ca i 2.52 -2.41 42.4 ... 26.8 43.2 27.7 17.6 33.6 19.8 21.9 20.8 ... 46.0 22.3 33.2 29.5 51.56330.09 Cr i 0.94 -2.92 127.0 63.7 141.3 127.1 146.4 127.9 141.2 141.4 127.0 132.2 36.7 169.3 116.1 143.9 140.8 146.26645.13 Eu ii 1.37 0.20 71.3 15.9 53.5 78.2 49.1 55.0 69.0 79.7 56.2 52.8 99.1 73.9 53.1 68.6 59.5 61.55369.96 Fe i 4.37 0.54 167.3 136.9 218.7 145.2 142.3 143.1 173.4 160.3 166.5 165.7 128.4 163.3 132.5 155.5 160.1 119.15383.37 Fe i 4.31 0.50 185.9 150.1 150.5 171.1 157.9 160.1 188.8 162.0 157.8 161.6 114.1 182.2 161.0 174.9 159.2 170.35386.34 Fe i 4.16 -1.74 53.4 ... ... ... 38.8 57.5 62.9 ... 44.6 51.8 ... 61.2 27.4 33.5 45.2 53.95393.17 Fe i 3.24 -0.92 199.6 180.7 198.8 217.9 209.6 233.1 216.2 200.3 218.7 212.3 176.0 239.2 187.7 207.6 221.0 214.65395.22 Fe i 4.45 -1.73 29.4 ... 29.8 30.2 23.3 30.1 31.4 29.3 23.1 ... ... 34.6 34.9 ... 30.1 31.45405.79 Fe i 0.99 -1.85 ... ... ... ... ... ... ... ... ... ... 262.0 ... ... ... ... ...5415.19 Fe i 4.39 0.51 155.8 145.2 151.4 174.2 167.5 186.8 152.5 155.3 176.2 167.1 108.5 168.1 162.7 166.5 174.0 184.95417.04 Fe i 4.42 -1.42 58.4 26.7 38.4 48.0 39.8 51.5 37.4 40.4 34.9 66.0 ... 53.7 28.8 52.4 35.7 71.35434.53 Fe i 1.01 -2.12 ... 299.9 ... ... ... ... ... ... ... ... 245.4 ... ... ... ... ...5436.30 Fe i 4.39 -1.35 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...5464.29 Fe i 4.14 -1.62 82.1 ... 59.5 90.4 53.5 ... 51.8 70.7 52.6 60.1 27.0 ... 49.2 52.6 62.9 73.55470.09 Fe i 4.45 -1.60 15.1 ... 28.2 24.7 25.2 17.5 25.4 24.2 33.7 21.2 16.8 ... 26.5 22.1 27.1 27.75501.48 Fe i 0.96 -3.05 251.8 238.7 266.5 267.8 275.4 254.1 274.3 260.6 257.5 270.8 185.7 274.6 242.7 273.0 251.4 286.45506.79 Fe i 0.99 -2.79 ... 234.7 ... ... ... ... ... ... ... ... 214.3 ... 279.6 ... ... ...5539.29 Fe i 3.64 -2.59 68.3 33.8 49.5 67.2 59.1 ... 57.5 58.8 58.1 59.5 ... 78.3 59.9 63.0 55.9 76.05586.77 Fe i 3.37 -0.10 235.9 177.8 206.8 244.4 239.6 245.6 243.9 225.7 228.8 224.0 143.4 248.2 227.3 245.4 228.6 250.26120.26 Fe i 0.91 -5.94 79.1 49.5 76.0 92.1 109.1 75.3 88.9 83.4 95.8 85.6 21.9 92.5 83.9 87.7 97.7 89.86136.62 Fe i 2.45 -1.50 ... 206.5 ... ... ... ... ... ... ... ... 217.6 ... ... ... ... ...6137.00 Fe i 2.20 -2.95 ... 129.5 ... ... ... ... ... ... ... ... ... ... ... ... ... ...6151.62 Fe i 2.18 -3.37 140.1 95.0 138.2 142.5 147.0 125.3 137.1 138.8 125.4 134.2 80.9 147.5 135.1 128.9 134.2 137.76157.75 Fe i 4.07 -1.26 133.4 50.1 105.8 125.1 100.1 104.2 113.4 110.4 104.8 115.1 82.2 151.9 86.0 104.4 104.6 126.46159.38 Fe i 4.61 -1.97 ... ... ... 23.5 26.4 ... ... 18.7 19.1 ... ... 65.3 ... ... 16.9 34.56165.36 Fe i 4.14 -1.47 83.7 32.6 67.6 64.3 64.9 67.5 58.8 65.4 74.7 68.5 24.2 92.0 63.1 72.8 71.0 75.56173.34 Fe i 2.22 -2.85 162.7 113.0 156.6 171.8 169.0 156.1 171.4 167.6 157.9 158.9 121.6 176.5 156.0 158.8 157.1 154.76180.20 Fe i 2.73 -2.78 133.2 101.8 138.6 119.5 149.9 116.9 123.1 139.7 154.7 132.3 59.3 85.7 132.5 142.0 143.8 147.36187.99 Fe i 3.94 -1.58 84.5 43.5 75.3 79.9 84.2 75.2 67.6 81.8 72.9 80.1 40.8 ... 63.7 86.5 76.2 69.46200.31 Fe i 2.61 -2.44 158.0 120.6 143.2 152.2 148.0 159.7 153.8 152.7 151.5 149.3 98.8 145.4 154.9 163.4 155.3 158.06213.43 Fe i 2.22 -2.66 210.8 147.0 194.4 200.8 184.6 189.0 192.7 183.8 181.6 211.8 143.0 199.2 194.4 183.9 187.4 187.56219.29 Fe i 2.20 -2.44 206.1 160.9 189.8 197.5 207.2 ... 198.5 184.7 204.1 202.2 133.7 212.0 191.6 202.5 189.9 201.26226.74 Fe i 3.88 -2.20 68.4 32.8 52.3 65.2 55.6 57.7 52.5 57.2 64.5 53.7 29.7 110.4 46.9 70.8 60.1 67.66252.57 Fe i 2.40 -1.76 205.2 180.6 207.3 199.2 199.8 198.6 210.1 205.0 194.4 205.9 153.3 223.5 194.1 203.8 199.8 202.16265.13 Fe i 2.18 -2.55 178.4 137.1 186.1 173.1 196.7 169.8 190.6 185.0 167.5 195.1 125.9 195.1 176.7 189.2 182.5 199.86271.28 Fe i 3.32 -2.96 65.9 17.8 52.9 57.7 64.5 54.3 52.2 68.3 71.0 51.7 21.1 57.0 64.9 53.0 67.1 77.76297.80 Fe i 2.22 -2.74 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6301.50 Fe i 3.65 -0.72 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6307.85 Fe i 3.64 -3.27 16.1 ... ... 24.0 ... ... ... ... ... ... ... ... ... ... ... 23.36322.69 Fe i 2.59 -2.43 162.5 127.2 153.4 164.7 164.6 162.2 164.7 160.1 160.6 172.0 111.3 178.3 158.5 153.4 160.5 179.36330.85 Fe i 4.73 -1.22 42.9 ... 31.3 48.0 34.2 41.5 37.2 43.4 31.8 27.0 18.2 48.2 39.8 42.7 39.1 37.66335.33 Fe i 2.20 -2.23 206.3 173.4 208.8 210.1 218.4 209.7 222.1 221.7 210.4 211.5 141.0 231.0 216.4 216.5 211.3 221.36336.82 Fe i 3.69 -1.05 152.0 126.2 141.9 143.0 139.4 138.7 144.1 144.5 170.1 141.9 98.9 163.1 135.3 141.4 136.1 163.96344.15 Fe i 2.43 -2.92 184.5 117.9 174.1 185.6 165.2 166.3 173.1 177.7 167.8 163.8 118.9 206.3 160.1 181.2 163.1 187.86355.04 Fe i 2.84 -2.29 165.2 116.7 170.7 164.0 159.6 157.0 165.8 163.6 154.9 168.7 87.9 172.1 143.3 172.9 166.7 174.96380.75 Fe i 4.19 -1.50 93.9 33.1 73.6 ... 77.9 80.3 77.1 79.7 73.4 65.7 28.2 121.3 99.8 56.4 82.7 86.66392.54 Fe i 2.28 -3.95 88.2 51.5 97.3 93.7 90.2 83.9 86.1 82.7 94.8 92.9 31.1 108.3 89.6 80.9 90.9 ...6393.61 Fe i 2.43 -1.63 251.5 205.5 239.4 231.0 255.6 245.3 244.3 235.1 244.5 235.4 169.2 261.1 228.0 255.6 236.7 259.46408.03 Fe i 3.69 -1.00 172.0 122.9 152.2 143.3 141.6 153.1 158.7 137.8 143.7 146.3 94.8 157.1 123.4 166.6 154.6 ...6419.96 Fe i 4.73 -0.24 112.5 73.9 98.2 104.6 110.7 94.8 113.4 121.4 110.9 108.3 54.0 135.0 99.1 113.0 102.4 106.66421.36 Fe i 2.28 -2.01 234.4 182.0 214.5 244.3 233.0 214.8 244.0 221.2 211.1 223.1 175.3 243.3 224.9 230.0 216.7 228.16430.86 Fe i 2.18 -1.95 261.9 201.4 250.8 269.5 272.7 264.8 269.7 244.4 251.7 245.5 180.9 297.6 240.2 251.8 245.5 262.06481.87 Fe i 2.27 -2.98 154.2 126.2 167.0 181.5 168.9 158.4 195.0 159.7 163.0 162.9 136.6 176.4 172.4 167.1 158.1 153.16498.94 Fe i 0.96 -4.69 182.7 131.5 190.8 204.4 196.2 189.8 186.7 188.1 186.6 185.1 134.8 212.6 185.1 183.0 180.5 218.16533.93 Fe i 4.55 -1.46 47.3 24.4 39.8 60.5 55.5 41.5 46.8 36.7 55.2 45.5 25.9 55.6 50.5 46.1 46.4 49.66556.81 Fe i 4.79 -1.72 18.6 ... ... 24.9 21.0 ... 21.1 ... 20.8 ... 20.7 ... 27.2 18.4 20.6 20.46569.22 Fe i 4.73 -0.42 134.3 61.1 98.3 118.4 112.2 84.9 106.9 109.0 103.9 110.9 89.4 136.6 117.0 105.7 103.6 110.46574.23 Fe i 0.99 -5.02 175.0 142.2 174.0 173.2 ... ... 177.1 ... 160.1 171.1 98.8 222.9 144.9 169.5 ... ...6593.88 Fe i 2.43 -2.39 231.2 180.8 195.1 192.9 225.9 205.3 193.8 210.1 182.7 192.4 90.7 246.4 151.8 211.4 227.0 203.76597.56 Fe i 4.79 -1.07 93.4 48.1 50.0 63.7 79.3 67.6 71.8 66.2 57.4 60.8 30.7 96.8 49.7 60.1 72.0 73.46608.03 Fe i 2.28 -3.94 94.3 58.4 107.6 116.4 104.8 85.5 99.0 91.1 96.0 97.5 60.1 101.3 98.9 97.7 112.8 97.56609.12 Fe i 2.56 -2.66 162.7 143.0 168.4 185.4 169.7 163.5 158.3 160.6 164.5 165.9 137.5 161.4 170.6 165.5 167.5 159.96627.54 Fe i 4.54 -1.68 39.8 ... 27.9 31.2 25.7 22.6 25.2 22.2 31.8 20.9 34.8 27.9 47.2 39.8 18.6 22.66633.76 Fe i 4.56 -0.82 122.6 53.0 103.2 123.2 95.0 109.5 94.9 112.9 112.6 ... 48.1 128.5 118.7 116.7 100.8 111.76646.93 Fe i 2.60 -3.99 56.4 34.6 63.6 69.7 55.1 49.0 52.6 56.0 59.1 50.2 ... 72.2 65.1 70.3 47.5 66.66653.85 Fe i 4.15 -2.52 33.6 ... ... 28.2 24.2 44.9 ... 30.2 34.7 39.6 34.3 37.1 ... 42.9 23.5 40.65414.08 Fe ii 3.22 -3.61 29.1 ... ... ... 19.6 19.0 39.2 21.9 27.3 43.1 25.4 ... 24.4 ... ... 25.55425.25 Fe ii 3.20 -3.36 52.3 37.0 41.5 52.5 40.0 49.2 38.1 44.9 69.4 54.5 50.3 36.7 51.9 45.2 36.7 70.56149.25 Fe ii 3.89 -2.72 44.7 25.4 31.8 ... 25.2 ... 40.8 37.5 36.1 33.3 33.6 49.8 24.7 28.9 29.1 ...6432.68 Fe ii 2.89 -3.71 59.3 57.8 54.5 55.9 50.9 60.5 55.6 59.2 53.4 53.1 59.5 49.9 54.3 64.5 50.9 51.36456.39 Fe ii 3.90 -2.08 78.7 74.6 51.0 50.9 ... 44.3 46.7 123.6 78.7 65.5 ... 68.0 27.9 84.7 ... ...6320.43 La ii 0.17 -1.56 86.9 23.0 74.1 119.7 56.6 72.0 80.4 83.5 65.2 71.8 87.8 119.3 61.4 73.2 71.9 100.76390.46 La ii 0.32 -1.40 74.4 31.0 74.9 80.7 56.3 56.1 65.2 76.9 62.8 84.1 76.2 85.3 61.3 81.2 74.7 69.75528.41 Mg i 4.35 -0.36 229.7 168.1 218.5 234.1 214.5 219.7 211.1 203.8 201.3 213.3 179.4 221.2 203.4 211.4 215.0 205.7



λ(Å) elem χ gf 113 115 123 125 132 135 138 140 141 146 147 148 149 150 151 1556318.72 Mg i 5.11 -1.97 ... 16.2 34.9 49.1 34.8 ... 35.6 46.7 39.3 30.9 18.4 49.4 ... 48.0 35.6 ...6319.24 Mg i 5.11 -2.21 ... 16.2 ... ... ... ... 20.5 ... ... 16.8 ... 31.6 29.6 15.2 25.0 43.26319.49 Mg i 5.11 -2.43 ... ... ... ... 17.4 ... ... 29.5 ... ... 30.2 ... ... ... ... ...5420.36 Mn i 2.14 -1.46 200.1 95.8 203.6 197.0 220.8 196.5 222.4 195.4 174.3 207.8 63.7 205.3 181.0 213.1 211.6 210.35432.55 Mn i 0.00 -3.80 263.5 179.6 254.6 266.5 291.7 262.1 ... 262.5 255.0 255.6 124.9 ... 241.8 293.8 261.2 294.85516.77 Mn i 2.18 -1.85 ... 47.3 147.1 174.4 162.9 148.4 175.0 136.1 144.6 138.2 32.4 162.0 ... 155.6 148.7 170.46154.23 Na i 2.10 -1.56 41.3 ... ... 20.5 20.4 20.9 ... ... 21.9 20.2 ... 46.1 ... 33.3 24.7 27.36160.75 Na i 2.10 -1.26 59.5 ... 31.9 66.9 22.8 49.0 ... 33.4 19.6 26.1 18.9 81.9 41.0 42.1 43.0 58.05416.38 Nd ii 0.86 -0.98 27.2 19.7 19.3 41.2 34.4 28.6 16.4 ... 27.1 38.4 16.3 50.2 ... 35.8 38.9 59.25431.54 Nd ii 1.12 -0.47 72.2 ... 48.1 72.8 58.9 63.0 54.4 58.8 51.9 57.6 62.1 103.5 ... 66.2 51.0 83.45485.71 Nd ii 1.26 -0.12 44.5 ... 37.1 47.6 33.0 37.2 36.3 33.1 32.1 41.4 58.7 50.5 29.4 32.4 36.1 26.95578.73 Ni i 1.68 -2.67 143.8 101.6 125.0 142.4 131.0 144.8 178.3 131.1 141.7 139.2 77.9 157.8 123.1 130.2 141.3 159.85587.87 Ni i 1.93 -2.37 178.7 110.3 160.3 167.9 174.8 168.4 160.2 154.8 155.5 164.5 77.3 168.8 146.9 165.8 164.0 161.15589.37 Ni i 3.90 -1.15 36.9 16.5 17.4 27.1 26.4 30.7 18.4 16.8 27.5 28.4 23.6 42.3 23.3 26.4 33.0 ...5593.75 Ni i 3.90 -0.79 39.5 24.1 32.6 32.6 30.3 36.5 50.2 40.1 38.7 47.1 21.4 56.9 22.9 30.6 34.2 57.46128.97 Ni i 1.68 -3.39 94.9 60.5 87.9 90.5 86.8 88.6 94.2 85.4 81.7 90.6 38.2 106.1 80.3 80.6 80.3 98.76130.14 Ni i 4.27 -0.98 ... ... ... 15.2 16.4 ... ... 27.6 19.6 ... ... 28.5 ... ... ... 18.36177.25 Ni i 1.83 -3.60 60.9 30.1 62.8 59.5 71.7 57.9 52.1 69.8 59.9 61.2 15.5 56.6 52.3 59.5 55.4 68.96186.72 Ni i 4.11 -0.90 33.4 ... 30.8 39.1 ... 18.6 28.9 20.7 20.9 18.1 ... 63.2 22.5 ... 24.1 46.16204.61 Ni i 4.09 -1.15 ... 17.6 ... 21.5 24.8 ... ... ... ... 23.7 ... ... ... ... 27.2 26.16223.99 Ni i 4.10 -0.97 26.4 ... 28.9 33.5 34.3 39.7 37.0 ... 32.9 26.1 ... 61.9 ... 36.7 37.2 34.16230.10 Ni i 4.11 -1.20 36.7 ... 29.3 ... ... 29.4 ... ... ... ... ... ... ... 30.0 ... 41.06322.17 Ni i 4.15 -1.21 ... ... 15.2 23.0 ... ... ... ... ... 15.4 ... ... 23.1 ... ... 40.06327.60 Ni i 1.68 -3.09 119.6 83.7 118.4 111.7 128.7 102.5 123.2 133.8 95.5 119.3 54.5 106.5 105.5 96.0 113.4 134.16378.26 Ni i 4.15 -0.82 33.2 ... 27.8 44.3 33.8 ... 32.9 35.4 32.5 23.6 ... 66.8 ... ... 27.1 ...6384.67 Ni i 4.15 -1.00 33.9 21.6 23.3 ... 18.5 33.5 31.6 23.5 33.3 16.9 16.8 46.2 22.6 36.5 ... 46.16482.80 Ni i 1.94 -2.85 101.1 66.4 107.9 114.1 126.7 119.3 111.3 116.9 98.8 108.5 62.4 97.0 95.9 105.9 114.7 115.16586.32 Ni i 1.95 -2.79 194.5 133.9 128.3 77.4 184.8 163.8 155.5 150.6 148.5 144.5 ... 244.2 45.6 171.5 190.1 198.36598.61 Ni i 4.24 -0.93 ... ... ... ... ... 32.1 21.3 30.7 22.5 ... ... 28.0 27.6 ... ... ...6635.14 Ni i 4.42 -0.75 ... ... 18.4 21.2 ... ... ... ... ... ... ... ... ... ... ... 20.36300.31 O i 0.00 -9.75 ... ... ... ... 60.9 ... ... ... ... ... ... ... ... ... 49.7 ...6363.79 O i 0.02 -10.25 25.3 ... 43.7 27.9 35.3 48.6 56.9 35.5 27.7 38.8 ... 76.2 29.8 30.0 42.6 46.15526.82 Sc ii 1.77 0.03 110.3 114.6 109.6 119.6 111.5 113.9 111.5 103.7 107.1 115.0 104.1 106.8 111.5 113.6 105.3 106.36245.62 Sc ii 1.51 -0.97 79.5 48.8 86.4 77.6 69.4 69.4 60.4 62.7 59.3 67.5 41.8 95.9 52.3 80.3 67.5 64.66309.90 Sc ii 1.50 -1.52 55.7 ... 42.8 46.8 ... ... 44.5 ... 27.2 ... ... 68.8 ... 41.6 ... ...6604.60 Sc ii 1.36 -1.31 105.8 63.6 99.6 97.4 82.0 93.5 78.6 83.5 73.7 90.6 80.8 109.0 90.8 77.5 92.9 98.36125.03 Si i 5.62 -1.57 ... ... ... ... 20.1 20.8 ... ... ... ... ... ... 31.7 ... ... ...6142.48 Si i 5.62 -1.51 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6145.02 Si i 5.61 -1.37 25.9 ... 19.3 15.4 ... ... ... ... 21.0 16.3 ... ... 21.8 25.1 ... ...6155.14 Si i 5.62 -0.80 58.8 ... 42.5 46.5 38.9 40.0 50.4 35.3 30.6 52.0 24.6 59.3 47.5 52.5 31.9 49.26237.33 Si i 5.61 -1.02 38.7 16.0 20.6 31.0 21.7 18.8 ... 20.6 ... 24.1 16.8 31.6 21.1 25.6 ... 24.16243.82 Si i 5.61 -1.27 ... ... ... 37.4 ... ... ... 22.8 ... 22.0 18.5 59.8 ... ... 29.3 ...5490.16 Ti i 1.46 -0.93 110.1 58.0 110.6 116.7 120.2 91.1 122.7 109.7 102.3 114.6 27.2 130.8 82.4 111.7 118.5 128.85503.90 Ti i 2.58 -0.19 67.9 37.5 58.1 69.5 68.6 61.1 78.3 64.8 48.9 52.4 17.1 101.7 31.1 60.6 65.1 98.46126.22 Ti i 1.07 -1.42 116.7 68.4 139.4 137.9 148.0 144.5 150.6 146.2 118.2 132.2 51.4 192.2 112.1 145.7 139.4 153.26220.50 Ti i 2.68 -0.14 ... ... ... 64.9 ... 76.8 ... ... ... ... 27.1 ... ... ... ... 91.76258.10 Ti i 1.44 -0.35 ... 98.9 190.7 ... 225.1 187.2 206.2 187.1 147.7 182.3 79.2 229.9 152.7 188.5 196.3 219.26303.77 Ti i 1.44 -1.57 93.2 42.4 80.8 107.9 91.6 91.4 104.1 86.3 79.1 85.1 ... 137.0 61.7 102.3 91.0 106.36312.24 Ti i 1.46 -1.55 ... ... ... 92.2 88.8 89.8 95.7 83.1 ... 71.3 15.5 ... 67.4 ... 86.2 93.46336.10 Ti i 1.44 -1.74 68.1 26.4 74.0 76.7 79.4 76.6 90.5 68.3 65.0 73.7 ... 104.9 ... 84.8 78.0 86.66508.12 Ti i 1.43 -2.05 38.7 ... 46.6 68.8 ... 61.0 61.7 52.8 30.8 44.5 ... 100.9 40.0 52.2 54.7 76.26556.08 Ti i 1.46 -1.07 ... 63.2 143.9 ... ... ... ... ... 111.7 ... ... ... ... 139.5 ... ...6599.13 Ti i 0.90 -2.09 131.4 67.1 134.4 152.6 164.2 126.2 147.0 141.4 94.4 127.9 48.6 208.9 99.1 132.2 146.6 157.76666.53 Ti i 1.46 -1.62 ... ... 16.4 34.4 ... 27.9 27.4 16.3 ... ... ... 39.6 ... 15.5 28.1 32.95418.77 Ti ii 1.58 -2.11 106.8 93.4 88.2 102.3 109.2 97.9 108.2 109.0 104.4 99.2 82.3 97.1 122.5 112.2 115.8 107.16219.94 Ti ii 2.06 -2.82 ... 27.4 ... ... ... ... ... ... 32.2 ... ... 60.9 ... ... 29.2 32.86559.58 Ti ii 2.05 -2.02 ... 36.1 61.0 ... ... ... ... ... 59.4 ... ... 57.0 ... 66.1 ... ...6606.95 Ti ii 2.06 -2.79 50.9 28.4 44.2 54.1 32.3 31.4 31.2 36.8 29.3 ... 61.6 37.6 50.0 33.7 35.9 28.76680.13 Ti ii 3.09 -1.86 33.2 18.7 ... 29.1 ... 21.6 22.9 ... 22.5 ... 24.3 ... 21.5 35.0 32.5 29.36119.53 V i 1.06 -0.32 82.3 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6128.33 V i 1.05 -2.30 ... ... ... ... ... 23.1 ... 17.2 ... ... ... 43.4 ... 22.1 ... 19.46135.37 V i 1.05 -0.75 100.8 25.1 100.6 108.0 123.9 100.8 105.6 95.6 93.4 88.5 ... 153.8 75.9 90.9 97.9 122.06150.15 V i 0.30 -1.79 119.9 59.5 143.2 142.7 168.2 130.2 166.0 141.7 108.0 136.4 ... 177.8 102.1 157.5 147.7 169.26199.19 V i 0.29 -1.29 171.2 95.8 171.8 180.1 194.0 165.2 199.6 182.2 152.7 161.8 36.4 207.4 139.5 181.2 187.1 213.46216.36 V i 0.28 -0.81 160.3 78.5 175.3 189.1 187.8 175.4 199.9 184.7 144.7 170.8 52.3 256.3 154.2 187.2 187.2 207.06224.51 V i 0.29 -2.01 120.1 ... 121.9 128.5 137.2 113.9 145.1 123.8 103.5 122.7 21.9 177.8 103.8 119.4 132.3 153.66233.20 V i 0.28 -2.07 79.7 22.3 99.8 111.8 105.5 76.6 96.4 83.2 72.7 80.9 ... 123.7 53.8 90.8 87.9 121.36243.11 V i 0.30 -0.98 ... 124.7 ... ... ... ... ... ... ... ... 58.6 ... ... ... ... ...6251.82 V i 0.29 -1.30 130.4 58.6 144.4 149.5 161.7 135.8 164.0 149.2 115.8 128.4 25.6 188.7 124.9 145.8 143.9 153.26274.66 V i 0.27 -1.67 ... ... ... ... 133.2 126.7 142.6 114.0 ... 113.2 19.0 ... 77.9 ... 122.1 148.46357.29 V i 1.85 -0.91 21.3 ... ... 36.8 30.5 24.3 ... 23.6 16.9 26.3 36.0 65.0 ... 28.3 ... 48.96452.32 V i 1.19 -1.21 95.2 32.5 73.5 96.5 103.1 86.9 90.8 79.7 61.2 67.2 ... 151.8 22.2 88.8 72.7 120.26504.19 V i 1.18 -1.23 55.7 ... 44.8 71.0 53.9 21.6 51.6 47.1 49.9 38.2 ... 77.6 34.0 57.2 41.9 47.96531.41 V i 1.22 -0.84 ... 29.1 66.0 ... ... ... ... ... 71.7 ... ... 100.9 ... 61.7 ... ...5402.78 Y ii 1.84 -0.51 40.7 ... 36.9 50.8 ... 34.7 32.2 48.4 22.9 37.6 22.9 45.7 43.3 ... ... 59.26362.35 Zn i 5.80 0.14 ... ... 31.9 ... ... ... ... ... 46.5 ... ... ... ... ... ... ...6127.48 Zr i 0.15 -1.06 68.5 19.4 76.8 86.5 89.3 80.7 87.0 95.8 52.4 68.5 18.5 109.5 64.6 75.7 79.8 104.16140.46 Zr i 0.52 -1.41 ... ... 18.8 21.7 ... ... 25.1 22.7 ... ... ... 53.4 ... ... ... 18.96143.18 Zr i 0.07 -1.10 83.7 ... 77.0 108.7 78.8 70.2 107.7 78.6 60.7 72.7 23.9 145.5 57.6 73.9 65.8 111.06192.95 Zr i 0.54 -2.07 ... ... 17.6 23.2 ... ... ... ... ... ... ... ... ... ... ... 18.1



Equivalent witdh, one star per column, BLxxxλ(Å) elem χ gf 156 158 160 163 166 168 171 173 180 185 190 195 196 197 203 2046141.73 Ba ii 0.70 -0.08 225.1 ... 240.2 286.5 266.8 229.1 238.0 254.3 ... 271.5 221.8 198.9 233.8 248.2 235.4 226.66496.91 Ba ii 0.60 -0.38 244.0 250.3 248.0 ... 254.1 238.6 221.7 265.0 290.5 279.9 230.2 229.5 231.6 294.5 245.7 293.46122.23 Ca i 1.89 -0.32 215.2 219.7 242.5 287.1 251.3 245.5 247.3 258.7 210.9 254.2 249.5 205.9 225.2 236.6 245.0 234.66156.03 Ca i 2.52 -2.39 ... 22.1 32.8 40.5 ... ... ... 29.0 30.0 44.3 20.5 ... 20.3 37.2 28.5 ...6161.30 Ca i 2.52 -1.27 82.4 117.5 119.1 152.0 138.7 132.7 115.0 143.1 128.1 146.0 122.7 93.3 111.4 140.4 143.2 114.86162.17 Ca i 1.90 -0.32 236.3 240.9 261.9 299.6 263.2 271.3 268.3 285.0 256.5 294.3 254.4 238.4 251.0 253.5 251.8 270.06166.44 Ca i 2.52 -1.14 89.8 102.1 122.8 137.7 117.9 110.2 121.7 132.9 108.8 132.6 134.0 93.0 115.0 127.3 128.5 112.06169.04 Ca i 2.52 -0.80 128.6 116.8 136.9 160.1 143.0 154.3 147.0 145.1 101.7 154.0 145.7 108.5 128.6 124.5 148.2 140.96169.56 Ca i 2.52 -0.48 131.8 148.2 162.7 181.1 174.7 173.7 177.3 156.9 127.1 165.2 155.1 130.8 156.6 145.7 161.1 156.56439.08 Ca i 2.52 0.39 200.2 228.1 212.2 251.3 229.0 232.1 243.3 232.7 ... 236.3 246.8 202.8 216.5 ... 237.1 ...6455.60 Ca i 2.52 -1.29 81.4 127.6 154.4 225.1 178.2 184.1 163.1 198.1 ... 171.9 166.4 124.8 110.0 ... 106.9 ...6471.67 Ca i 2.52 -0.76 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6493.79 Ca i 2.52 -0.32 168.5 163.9 179.8 198.6 192.5 186.1 199.3 193.7 185.8 182.5 192.0 152.5 182.4 207.2 192.7 208.76499.65 Ca i 2.52 -0.82 109.3 117.4 123.9 150.4 124.5 133.7 131.7 139.9 132.3 142.3 130.2 106.8 124.4 151.7 137.3 132.16508.84 Ca i 2.52 -2.41 ... 16.5 32.7 42.6 23.9 24.2 26.2 20.6 24.7 45.6 25.6 18.7 21.1 43.6 38.5 ...6330.09 Cr i 0.94 -2.92 104.5 117.7 122.9 152.9 117.2 131.3 150.5 142.5 120.8 158.1 133.9 92.1 137.3 136.1 132.3 126.36645.13 Eu ii 1.37 0.20 60.6 73.1 60.1 79.0 56.8 67.5 57.4 61.5 73.6 69.8 61.4 50.0 58.8 76.5 68.6 103.25369.96 Fe i 4.37 0.54 162.0 166.7 206.6 171.3 151.9 168.6 178.3 143.9 155.9 ... 160.8 192.0 199.2 167.9 147.9 153.25383.37 Fe i 4.31 0.50 169.9 166.2 139.3 175.2 185.8 173.5 171.8 169.9 150.5 163.7 171.9 157.3 163.1 172.5 140.9 156.95386.34 Fe i 4.16 -1.74 53.6 53.1 54.8 47.3 38.5 ... 54.9 56.4 54.0 67.3 62.0 25.5 44.7 37.7 48.4 58.95393.17 Fe i 3.24 -0.92 183.9 210.6 219.9 230.0 229.4 216.6 211.2 251.7 193.1 218.9 204.4 181.8 194.5 179.7 220.3 197.75395.22 Fe i 4.45 -1.73 20.7 ... ... 34.8 35.1 22.3 16.7 24.3 ... 26.4 42.7 ... 19.6 21.3 ... ...5405.79 Fe i 0.99 -1.85 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...5415.19 Fe i 4.39 0.51 136.8 151.1 159.0 193.5 181.7 175.9 195.7 191.2 137.7 163.0 176.3 163.9 169.2 150.8 174.7 153.35417.04 Fe i 4.42 -1.42 47.1 40.0 70.3 49.3 52.3 53.5 43.0 57.7 ... 53.1 48.2 46.1 46.2 ... 19.9 33.85434.53 Fe i 1.01 -2.12 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...5436.30 Fe i 4.39 -1.35 ... 101.2 ... ... ... ... ... ... ... ... ... ... ... ... ... ...5464.29 Fe i 4.14 -1.62 27.8 57.0 76.3 78.1 67.3 56.4 50.9 60.2 72.8 67.2 57.1 61.1 38.4 72.1 ... 54.15470.09 Fe i 4.45 -1.60 ... 25.4 28.3 42.6 22.5 16.5 ... 40.6 38.9 32.0 21.5 21.8 19.7 30.1 18.4 17.65501.48 Fe i 0.96 -3.05 234.0 233.3 257.0 278.4 266.8 264.8 282.2 272.6 232.1 255.2 280.6 243.2 262.5 274.4 261.9 230.65506.79 Fe i 0.99 -2.79 262.4 295.6 ... ... ... ... ... ... 290.0 ... ... 274.6 ... ... ... 295.95539.29 Fe i 3.64 -2.59 32.4 64.9 65.7 76.2 70.3 57.9 57.3 85.4 43.4 75.9 57.1 53.5 46.2 39.2 57.9 40.85586.77 Fe i 3.37 -0.10 227.9 246.8 223.8 250.6 246.9 237.1 230.5 256.1 168.4 232.4 238.7 221.8 220.4 202.2 243.4 195.16120.26 Fe i 0.91 -5.94 76.0 79.1 96.3 110.2 103.0 109.5 94.5 82.0 63.8 83.3 107.0 66.7 88.7 90.0 98.9 82.26136.62 Fe i 2.45 -1.50 ... ... ... ... ... ... ... ... ... ... ... 227.1 261.7 281.1 ... ...6137.00 Fe i 2.20 -2.95 ... ... ... ... ... ... ... ... ... ... ... 171.3 158.0 170.5 ... ...6151.62 Fe i 2.18 -3.37 113.5 140.1 130.2 133.1 138.0 122.9 129.2 143.7 122.6 144.0 143.0 126.9 128.2 129.0 133.4 126.16157.75 Fe i 4.07 -1.26 94.5 101.4 116.5 120.5 112.6 120.9 109.4 112.1 141.1 118.0 108.9 95.3 105.5 121.0 115.0 121.76159.38 Fe i 4.61 -1.97 20.1 ... ... 28.6 ... ... ... 20.8 ... 38.7 17.0 ... ... 24.3 20.3 18.56165.36 Fe i 4.14 -1.47 55.7 67.6 71.0 76.0 72.6 67.6 64.9 69.3 66.4 87.0 80.6 66.7 67.5 69.9 73.3 54.56173.34 Fe i 2.22 -2.85 142.5 142.6 173.1 169.8 155.6 157.9 153.4 165.3 160.6 162.6 182.1 155.7 154.8 163.1 163.9 162.86180.20 Fe i 2.73 -2.78 111.3 119.3 126.8 ... 154.7 170.7 147.2 144.3 74.0 126.0 151.5 118.7 ... 102.6 137.8 133.26187.99 Fe i 3.94 -1.58 64.3 64.9 82.2 85.5 83.3 79.1 68.8 82.8 73.2 81.0 79.5 64.7 77.8 81.6 74.0 71.16200.31 Fe i 2.61 -2.44 145.8 135.5 155.6 168.8 148.7 150.1 155.4 157.8 147.7 151.6 159.6 132.3 148.6 130.0 156.1 145.86213.43 Fe i 2.22 -2.66 171.7 193.8 185.5 195.0 192.3 181.1 201.0 191.1 187.4 203.5 205.5 164.8 190.2 193.6 184.3 190.56219.29 Fe i 2.20 -2.44 187.0 207.9 201.6 214.2 204.7 187.9 193.5 211.4 192.1 223.9 205.5 183.7 193.6 202.9 198.7 201.76226.74 Fe i 3.88 -2.20 33.5 58.5 59.9 71.2 59.1 59.8 54.7 45.7 58.4 60.0 72.2 68.3 55.5 60.3 72.0 68.66252.57 Fe i 2.40 -1.76 188.6 190.2 210.8 222.0 217.8 200.5 202.5 209.8 198.5 199.5 223.8 194.3 211.5 214.7 190.6 202.66265.13 Fe i 2.18 -2.55 174.0 168.5 192.2 193.1 192.4 185.9 188.1 170.4 173.1 185.7 196.2 162.4 184.3 187.6 176.2 182.36271.28 Fe i 3.32 -2.96 47.9 52.8 55.0 66.9 60.2 61.6 53.1 72.0 61.2 62.3 61.9 44.7 51.7 61.7 55.3 70.66297.80 Fe i 2.22 -2.74 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6301.50 Fe i 3.65 -0.72 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6307.85 Fe i 3.64 -3.27 ... ... ... 23.2 ... ... ... 33.5 ... ... ... ... ... 18.3 ... ...6322.69 Fe i 2.59 -2.43 151.4 148.5 160.4 165.2 147.5 166.0 169.9 171.5 137.3 162.4 173.7 140.8 156.4 158.7 140.2 153.56330.85 Fe i 4.73 -1.22 22.6 31.5 39.7 38.0 38.9 36.2 26.2 38.3 35.7 43.1 37.8 25.8 25.2 34.7 28.5 50.56335.33 Fe i 2.20 -2.23 198.4 198.9 207.1 227.9 206.1 194.0 209.6 229.7 201.3 223.0 226.4 188.2 220.2 221.9 213.3 212.86336.82 Fe i 3.69 -1.05 153.7 138.2 142.5 153.3 138.0 154.9 144.5 164.3 144.8 144.1 154.6 137.6 149.2 144.0 134.1 140.96344.15 Fe i 2.43 -2.92 161.7 167.4 163.1 201.9 177.3 165.8 169.7 188.9 175.8 198.2 183.3 126.5 ... 163.6 168.5 160.76355.04 Fe i 2.84 -2.29 131.6 154.5 151.7 158.8 158.5 147.6 153.7 165.2 165.5 174.2 160.0 139.9 163.1 161.3 155.1 138.56380.75 Fe i 4.19 -1.50 50.2 88.4 74.7 93.3 67.6 80.9 68.7 96.5 72.3 99.7 90.6 77.2 80.1 75.6 84.5 99.06392.54 Fe i 2.28 -3.95 82.1 91.4 87.2 90.5 ... 76.7 88.9 116.1 75.4 93.4 97.6 66.6 88.9 91.1 98.1 94.16393.61 Fe i 2.43 -1.63 225.5 226.0 241.8 257.2 227.9 243.4 251.5 249.0 236.4 258.3 256.2 225.2 236.9 220.9 242.8 235.46408.03 Fe i 3.69 -1.00 124.1 126.1 145.9 177.4 151.3 144.5 165.7 152.3 ... 169.4 150.4 138.7 159.9 ... 128.7 ...6419.96 Fe i 4.73 -0.24 91.4 112.4 106.4 123.4 89.2 96.0 112.2 120.5 ... 113.2 115.1 104.2 103.8 ... 107.7 ...6421.36 Fe i 2.28 -2.01 209.8 226.0 215.6 242.9 231.2 194.9 223.8 255.2 ... 232.8 241.4 185.8 220.4 ... 222.5 ...6430.86 Fe i 2.18 -1.95 230.1 248.4 236.8 275.6 260.4 255.1 265.1 282.6 ... 261.4 279.5 227.1 232.2 ... 256.6 ...6481.87 Fe i 2.27 -2.98 147.4 148.1 176.7 169.4 183.8 160.2 175.7 161.7 172.7 185.0 171.9 151.2 158.9 178.4 176.9 197.56498.94 Fe i 0.96 -4.69 170.8 182.5 181.7 205.0 205.6 168.4 196.6 196.1 188.4 207.3 199.9 160.8 175.8 210.1 185.3 209.36533.93 Fe i 4.55 -1.46 29.3 48.7 61.8 56.6 41.8 38.0 44.7 50.9 ... 47.6 52.4 28.8 51.7 ... 42.2 ...6556.81 Fe i 4.79 -1.72 ... ... ... 15.9 24.1 ... ... ... 15.1 ... 19.7 22.0 ... 27.0 ... ...6569.22 Fe i 4.73 -0.42 90.2 107.9 112.2 136.1 109.8 104.0 107.7 106.7 117.5 126.6 108.0 93.8 98.3 121.4 118.0 154.96574.23 Fe i 0.99 -5.02 ... ... 167.7 ... 188.2 182.6 180.5 ... ... 187.1 197.2 ... 187.4 ... 168.0 ...6593.88 Fe i 2.43 -2.39 145.7 152.4 195.7 250.6 196.8 228.2 219.0 224.9 ... 208.3 228.4 180.0 196.7 ... 152.3 ...6597.56 Fe i 4.79 -1.07 32.5 51.8 55.0 87.4 70.8 74.9 81.6 65.7 ... 57.5 77.0 52.0 57.1 ... 50.1 ...6608.03 Fe i 2.28 -3.94 85.8 100.6 89.7 115.7 96.7 95.5 106.8 106.3 123.3 99.0 98.6 79.3 108.7 134.5 109.1 164.26609.12 Fe i 2.56 -2.66 165.1 162.7 150.2 190.8 166.6 150.6 171.3 184.7 184.2 161.9 170.8 159.1 159.7 185.0 173.3 225.56627.54 Fe i 4.54 -1.68 20.0 33.5 26.7 25.1 24.2 33.4 32.3 29.9 60.6 ... 33.8 20.8 23.3 49.4 41.7 110.86633.76 Fe i 4.56 -0.82 85.7 111.2 115.3 118.6 111.0 94.5 98.3 122.9 133.9 121.4 112.9 97.3 89.6 137.6 119.3 187.16646.93 Fe i 2.60 -3.99 38.6 59.8 ... 65.3 59.4 73.3 61.4 67.3 ... 64.6 65.2 36.6 62.4 72.8 55.3 115.76653.85 Fe i 4.15 -2.52 ... 40.3 ... 38.4 31.6 ... 34.8 29.7 36.2 32.1 27.3 21.9 ... 33.3 33.4 ...5414.08 Fe ii 3.22 -3.61 23.4 39.7 32.5 25.7 51.9 ... 36.1 27.0 ... ... 31.1 40.8 30.7 ... 29.3 ...5425.25 Fe ii 3.20 -3.36 50.8 45.4 42.3 49.5 38.5 61.4 56.9 34.6 53.4 44.7 48.6 48.0 28.5 35.0 57.8 41.76149.25 Fe ii 3.89 -2.72 41.4 35.9 29.4 25.8 ... 27.5 ... 25.8 ... 42.8 21.1 ... 36.5 29.5 ... 43.96432.68 Fe ii 2.89 -3.71 59.9 70.8 55.4 66.9 45.2 35.2 51.1 48.7 ... 54.0 54.5 66.5 45.2 ... 53.5 ...6456.39 Fe ii 3.90 -2.08 113.0 ... 69.2 ... 112.6 91.3 99.5 67.9 ... 43.3 ... 60.7 70.7 ... ... ...6320.43 La ii 0.17 -1.56 66.6 95.5 80.2 99.9 88.8 82.1 70.1 90.0 108.7 111.8 41.4 ... 75.6 79.3 73.3 53.46390.46 La ii 0.32 -1.40 73.9 70.9 ... 92.7 61.5 76.2 63.4 84.2 93.7 81.7 59.7 37.8 58.1 98.7 76.9 77.35528.41 Mg i 4.35 -0.36 202.6 214.5 206.5 221.8 216.6 205.1 214.3 225.0 209.3 228.0 209.7 196.9 208.5 209.9 212.6 217.3



λ(Å) elem χ gf 156 158 160 163 166 168 171 173 180 185 190 195 196 197 203 2046318.72 Mg i 5.11 -1.97 ... 37.4 29.8 41.1 37.8 32.0 38.7 36.9 33.1 47.8 29.6 32.7 ... 40.3 37.3 ...6319.24 Mg i 5.11 -2.21 20.3 35.0 ... ... 25.3 ... 30.1 25.5 ... 25.5 ... ... ... ... 28.9 ...6319.49 Mg i 5.11 -2.43 ... ... ... ... ... ... ... ... ... ... ... ... ... 31.2 ... ...5420.36 Mn i 2.14 -1.46 150.2 206.1 203.1 221.2 215.5 189.3 222.0 210.7 181.6 211.2 208.7 169.8 168.1 198.9 217.9 180.95432.55 Mn i 0.00 -3.80 214.4 255.9 264.5 ... 269.7 274.6 ... ... 241.8 272.8 287.4 191.7 258.7 297.8 297.4 254.95516.77 Mn i 2.18 -1.85 95.4 ... 151.6 166.5 151.8 133.5 138.1 ... 130.4 182.3 145.0 108.8 133.6 159.5 152.5 142.76154.23 Na i 2.10 -1.56 ... 22.3 ... 41.9 19.6 ... ... 25.5 19.3 49.6 ... 18.2 ... 23.6 24.3 ...6160.75 Na i 2.10 -1.26 ... 42.3 40.7 53.0 47.3 47.9 ... 59.0 45.5 54.0 ... 23.7 19.0 42.0 53.5 37.65416.38 Nd ii 0.86 -0.98 35.7 35.1 43.2 54.3 38.9 29.1 30.3 28.6 ... ... ... 16.6 28.8 ... ... ...5431.54 Nd ii 1.12 -0.47 ... 80.1 52.8 76.6 60.0 47.6 45.2 77.4 66.9 68.0 54.2 31.0 34.9 49.7 65.3 46.75485.71 Nd ii 1.26 -0.12 19.8 33.2 42.0 53.0 43.5 35.4 17.7 30.8 61.9 51.3 26.4 18.0 40.2 43.7 40.8 ...5578.73 Ni i 1.68 -2.67 112.7 147.7 142.8 145.1 147.2 142.7 128.4 156.1 130.1 150.4 143.3 115.8 132.9 135.1 136.1 121.85587.87 Ni i 1.93 -2.37 136.7 147.5 163.4 191.0 175.0 167.2 154.0 169.1 146.5 176.9 161.7 153.4 150.3 143.2 153.9 137.35589.37 Ni i 3.90 -1.15 ... ... 26.9 ... 19.9 23.5 31.0 36.2 25.2 45.0 19.1 22.3 21.7 35.6 23.8 32.35593.75 Ni i 3.90 -0.79 23.7 49.4 47.4 45.0 43.7 34.6 41.7 54.7 34.6 40.8 40.3 39.4 27.0 32.1 40.5 36.36128.97 Ni i 1.68 -3.39 68.1 79.7 76.2 84.7 82.6 89.9 84.1 83.6 79.3 95.2 87.8 69.8 72.5 86.7 98.0 76.86130.14 Ni i 4.27 -0.98 ... ... 16.4 23.1 ... ... ... ... ... ... ... ... ... 29.0 ... ...6177.25 Ni i 1.83 -3.60 54.1 53.2 69.2 79.8 51.3 62.1 69.6 55.4 44.1 64.2 67.1 38.7 55.5 57.8 56.6 60.86186.72 Ni i 4.11 -0.90 18.1 20.5 35.0 39.3 28.1 42.4 24.0 33.1 23.7 32.2 35.7 23.5 16.5 23.8 32.3 20.96204.61 Ni i 4.09 -1.15 ... 20.4 36.3 38.5 15.0 21.8 ... 25.7 ... 16.3 30.5 ... ... ... ... ...6223.99 Ni i 4.10 -0.97 ... ... 20.6 47.2 ... 22.3 ... 33.9 18.6 41.4 ... ... 21.7 36.9 38.3 ...6230.10 Ni i 4.11 -1.20 ... ... 29.2 ... ... ... 27.6 29.1 37.6 47.3 ... ... ... ... 20.5 ...6322.17 Ni i 4.15 -1.21 ... 30.6 ... ... ... 32.4 20.6 46.5 ... 26.7 ... ... ... ... ... ...6327.60 Ni i 1.68 -3.09 106.8 117.3 114.2 132.6 112.0 106.7 105.5 130.8 109.1 123.6 114.1 99.2 116.8 122.5 93.3 107.26378.26 Ni i 4.15 -0.82 20.6 33.7 28.8 33.4 44.1 27.9 38.4 ... 32.6 23.7 15.5 ... 27.7 ... 34.5 ...6384.67 Ni i 4.15 -1.00 ... ... 26.3 41.4 31.7 26.6 28.4 ... ... 42.2 ... ... ... 22.3 ... ...6482.80 Ni i 1.94 -2.85 105.1 116.7 109.1 120.5 106.3 97.9 101.4 113.7 104.0 113.9 106.0 96.5 102.1 131.1 111.4 109.46586.32 Ni i 1.95 -2.79 57.5 84.8 133.6 210.7 143.3 217.8 183.8 192.6 ... 176.2 185.4 109.3 179.5 ... 92.3 ...6598.61 Ni i 4.24 -0.93 ... ... 17.0 ... 30.3 ... 34.6 ... ... ... ... ... ... ... ... ...6635.14 Ni i 4.42 -0.75 29.3 ... ... ... ... ... ... ... 25.1 ... ... ... ... 38.8 25.7 ...6300.31 O i 0.00 -9.75 ... 66.7 ... ... ... ... ... ... ... ... ... ... ... ... ... ...6363.79 O i 0.02 -10.25 ... 43.1 ... 46.5 32.7 33.3 41.0 82.8 33.3 43.3 35.4 22.6 ... 42.8 34.7 38.55526.82 Sc ii 1.77 0.03 116.2 104.1 119.5 105.9 110.4 99.2 130.1 95.1 123.1 118.7 99.5 99.6 124.8 115.6 129.0 112.16245.62 Sc ii 1.51 -0.97 74.3 73.9 75.9 61.7 65.2 65.3 78.0 72.4 82.5 64.3 ... 62.0 71.6 57.8 82.2 65.56309.90 Sc ii 1.50 -1.52 ... ... 38.9 ... 41.7 41.5 39.1 ... 40.4 ... 29.6 ... 30.4 ... 47.4 46.86604.60 Sc ii 1.36 -1.31 79.4 96.7 80.9 105.7 58.9 84.3 102.2 88.2 101.6 74.8 76.6 89.8 81.3 112.3 101.7 130.06125.03 Si i 5.62 -1.57 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6142.48 Si i 5.62 -1.51 ... ... ... ... ... ... ... ... ... 18.4 ... ... ... ... ... ...6145.02 Si i 5.61 -1.37 20.6 17.7 21.8 ... ... ... ... ... ... 22.4 ... ... ... ... ... ...6155.14 Si i 5.62 -0.80 ... 38.4 36.0 45.6 29.4 32.6 ... 55.7 54.0 53.8 35.9 46.7 38.5 44.0 34.2 47.16237.33 Si i 5.61 -1.02 18.4 ... 31.0 26.0 ... 17.4 ... ... 31.4 ... 15.2 15.8 ... 24.4 19.9 ...6243.82 Si i 5.61 -1.27 15.7 31.6 33.4 ... 36.7 33.6 ... ... ... 43.3 ... 23.6 ... ... 43.2 ...5490.16 Ti i 1.46 -0.93 70.3 99.8 124.8 133.6 109.2 111.7 104.3 130.2 84.2 114.5 118.2 68.1 88.1 108.5 126.7 111.75503.90 Ti i 2.58 -0.19 23.3 60.9 71.0 87.9 64.4 66.8 51.1 76.5 45.1 73.1 68.1 25.3 47.4 58.7 62.7 41.76126.22 Ti i 1.07 -1.42 113.8 119.7 138.8 153.9 132.6 132.0 132.4 152.3 124.4 149.9 147.5 101.3 120.6 145.1 154.4 126.16220.50 Ti i 2.68 -0.14 ... 80.5 ... ... 69.0 ... ... 72.4 ... 90.4 ... ... ... ... ... ...6258.10 Ti i 1.44 -0.35 138.9 157.8 186.4 219.4 176.7 206.5 ... 207.7 149.3 197.1 191.4 117.3 158.3 205.0 185.6 178.46303.77 Ti i 1.44 -1.57 33.8 60.0 ... 86.3 87.1 98.3 71.8 105.4 78.6 105.8 79.9 44.6 79.1 90.3 80.7 78.96312.24 Ti i 1.46 -1.55 37.9 60.4 81.1 97.8 72.4 ... 68.2 88.5 ... 87.9 84.4 44.5 ... 89.6 89.7 75.16336.10 Ti i 1.44 -1.74 36.1 54.6 73.5 102.5 75.8 68.5 83.5 100.7 54.6 97.4 64.3 36.0 62.0 79.5 73.6 82.26508.12 Ti i 1.43 -2.05 ... 24.8 62.3 ... 56.1 42.3 42.8 59.6 33.9 64.9 45.2 22.9 33.8 65.3 54.9 57.56556.08 Ti i 1.46 -1.07 ... ... ... ... ... 133.3 ... ... 126.7 ... ... ... 123.4 ... ... ...6599.13 Ti i 0.90 -2.09 72.4 93.8 141.2 166.2 117.3 136.8 144.5 178.2 ... 154.9 138.3 75.9 119.3 ... 115.4 ...6666.53 Ti i 1.46 -1.62 ... 29.9 ... 19.5 ... 18.8 ... 22.5 24.3 42.7 ... 16.2 ... 30.3 20.7 ...5418.77 Ti ii 1.58 -2.11 95.7 108.8 110.7 113.7 102.1 112.0 124.6 113.1 89.6 101.4 107.3 94.0 94.2 104.3 100.3 97.96219.94 Ti ii 2.06 -2.82 ... 36.6 ... ... ... 24.0 ... 29.5 ... 31.7 29.9 19.0 20.8 ... ... ...6559.58 Ti ii 2.05 -2.02 ... ... ... ... ... ... ... ... 66.8 ... ... ... 50.8 ... ... ...6606.95 Ti ii 2.06 -2.79 36.6 43.4 43.2 59.3 41.4 21.7 30.1 49.3 62.8 41.2 ... 41.6 34.3 61.9 ... 116.76680.13 Ti ii 3.09 -1.86 ... 30.1 ... 41.2 35.1 ... 33.8 46.0 27.6 18.9 24.0 27.8 23.6 42.1 29.4 41.76119.53 V i 1.06 -0.32 ... ... ... 95.2 ... 85.0 83.8 ... ... ... 71.8 ... ... ... ... ...6128.33 V i 1.05 -2.30 ... ... ... 28.4 ... ... ... ... ... ... ... ... ... ... ... ...6135.37 V i 1.05 -0.75 59.0 74.2 98.4 125.6 94.6 92.1 90.6 120.3 88.8 120.2 98.7 44.8 76.0 116.9 107.6 79.76150.15 V i 0.30 -1.79 96.1 122.6 143.3 167.4 136.1 150.2 144.5 172.1 95.1 165.9 131.1 69.0 123.5 153.3 143.5 140.16199.19 V i 0.29 -1.29 111.7 150.8 204.2 202.9 166.0 178.2 177.3 214.4 144.4 195.4 186.9 101.1 155.0 180.9 193.7 ...6216.36 V i 0.28 -0.81 128.8 168.3 182.5 194.6 176.4 177.0 164.9 199.8 166.2 203.9 184.0 119.0 150.0 196.7 182.8 178.46224.51 V i 0.29 -2.01 73.1 101.4 134.9 147.3 121.0 142.6 130.0 147.7 101.6 150.4 120.1 72.8 99.1 136.0 136.8 108.56233.20 V i 0.28 -2.07 39.5 73.3 100.9 96.7 81.9 93.7 88.3 100.8 66.8 108.1 91.8 29.3 72.9 104.0 100.5 72.96243.11 V i 0.30 -0.98 ... 234.4 ... ... ... 284.4 ... ... ... ... ... 177.1 ... ... ... ...6251.82 V i 0.29 -1.30 97.9 116.9 144.0 149.8 144.7 144.5 135.5 158.8 128.1 172.4 147.0 83.6 130.8 161.6 151.6 134.56274.66 V i 0.27 -1.67 65.1 108.5 133.7 132.8 ... ... ... 134.1 ... 127.4 ... 57.2 ... 128.9 119.2 ...6357.29 V i 1.85 -0.91 ... 23.3 25.5 30.9 31.4 ... ... 35.0 ... 35.2 20.4 19.6 ... ... 26.4 ...6452.32 V i 1.19 -1.21 18.9 50.7 102.7 124.9 86.3 99.2 97.0 115.4 ... 110.9 105.5 32.6 64.3 ... 66.7 ...6504.19 V i 1.18 -1.23 18.6 35.6 56.3 68.4 48.3 47.8 58.9 53.9 44.1 42.7 47.6 15.5 33.7 64.6 54.6 46.06531.41 V i 1.22 -0.84 ... ... ... ... ... 80.0 ... ... ... ... ... ... 52.5 ... ... ...5402.78 Y ii 1.84 -0.51 ... 34.8 40.5 51.6 43.8 27.1 28.8 35.0 58.2 47.6 21.6 ... 36.3 53.9 30.6 26.66362.35 Zn i 5.80 0.14 ... ... ... ... ... 22.5 ... ... ... ... ... ... 17.2 ... ... ...6127.48 Zr i 0.15 -1.06 ... 75.5 77.0 104.5 77.4 81.4 76.3 101.7 65.5 88.5 72.5 ... 69.4 87.8 73.8 ...6140.46 Zr i 0.52 -1.41 ... ... 16.7 ... 19.2 ... ... ... 19.8 27.5 ... ... ... 26.2 27.3 ...6143.18 Zr i 0.07 -1.10 24.9 66.9 81.3 90.2 79.4 84.7 71.4 108.8 81.7 105.4 70.4 26.9 55.3 102.3 85.2 54.16192.95 Zr i 0.54 -2.07 ... ... ... ... ... ... ... 23.7 23.8 ... ... ... ... 19.6 ... ...



Equivalent witdh, one star per column, BLxxxλ(Å) elem χ gf 205 208 210 211 213 216 218 221 227 228 229 233 239 242 247 2506141.73 Ba ii 0.70 -0.08 267.8 259.2 265.1 271.8 252.9 282.6 ... 251.0 244.0 215.3 275.9 261.5 228.5 226.7 250.1 ...6496.91 Ba ii 0.60 -0.38 243.1 270.3 284.0 297.9 235.1 281.4 280.7 256.7 268.8 232.7 284.4 247.9 258.8 250.3 246.8 ...6122.23 Ca i 1.89 -0.32 244.8 250.7 233.3 279.0 246.7 276.6 ... 247.1 230.7 248.7 267.5 242.6 228.3 213.3 250.5 ...6156.03 Ca i 2.52 -2.39 ... 26.0 25.6 36.6 24.6 42.3 39.8 25.7 26.1 35.5 42.9 26.5 20.7 28.5 36.9 38.66161.30 Ca i 2.52 -1.27 150.3 142.6 145.4 155.4 137.5 160.3 163.3 127.7 136.1 129.6 143.4 141.2 108.0 106.0 145.6 167.06162.17 Ca i 1.90 -0.32 261.5 273.6 274.5 ... 268.1 ... ... 255.5 266.5 269.6 289.8 285.5 249.2 237.1 279.7 ...6166.44 Ca i 2.52 -1.14 133.0 132.0 119.9 146.1 117.2 136.1 148.2 117.9 117.6 123.8 158.7 121.9 118.0 96.5 133.1 174.56169.04 Ca i 2.52 -0.80 144.1 155.5 133.2 156.5 152.7 149.2 148.9 144.0 126.5 150.0 146.1 147.2 150.0 95.7 147.2 170.96169.56 Ca i 2.52 -0.48 164.3 170.2 176.8 167.0 159.8 171.8 183.3 167.8 169.8 153.2 176.3 164.3 146.0 117.5 169.5 210.36439.08 Ca i 2.52 0.39 225.4 228.8 231.7 234.7 227.2 ... 243.2 205.2 210.7 231.2 235.4 235.6 224.7 200.1 237.1 243.96455.60 Ca i 2.52 -1.29 192.1 161.5 184.7 113.3 117.5 ... ... 57.1 102.5 133.8 175.2 185.3 175.1 123.0 163.5 222.96471.67 Ca i 2.52 -0.76 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6493.79 Ca i 2.52 -0.32 195.2 181.3 191.7 191.3 182.7 222.4 223.7 179.6 178.8 189.8 208.7 196.6 180.4 160.4 163.1 190.36499.65 Ca i 2.52 -0.82 111.5 122.6 135.9 140.6 124.9 153.6 145.6 115.7 124.0 116.1 137.1 126.5 126.6 125.1 121.6 159.16508.84 Ca i 2.52 -2.41 31.7 36.6 32.1 50.9 25.0 67.1 58.0 36.9 24.0 18.6 41.6 33.1 22.5 16.0 25.5 67.46330.09 Cr i 0.94 -2.92 122.0 122.3 132.8 160.3 130.2 161.8 160.6 133.3 138.7 150.4 142.5 132.2 116.1 120.3 147.0 178.86645.13 Eu ii 1.37 0.20 57.2 52.8 72.2 74.0 68.7 86.3 68.4 83.5 76.2 50.7 80.2 65.1 40.1 47.2 66.4 62.55369.96 Fe i 4.37 0.54 167.4 163.5 164.9 168.3 147.9 175.4 146.3 201.7 148.9 217.2 ... 152.5 148.5 134.5 194.3 199.05383.37 Fe i 4.31 0.50 181.5 162.1 170.0 164.5 169.7 172.1 204.0 171.3 166.6 166.8 165.7 169.5 145.3 155.0 182.6 177.15386.34 Fe i 4.16 -1.74 60.5 52.5 53.0 ... 53.7 55.1 59.5 ... 22.1 46.8 51.8 61.5 72.8 25.3 60.5 ...5393.17 Fe i 3.24 -0.92 221.4 227.6 200.5 209.7 222.4 226.2 211.3 211.4 210.9 209.1 231.5 230.0 224.4 201.2 222.3 246.85395.22 Fe i 4.45 -1.73 25.6 24.3 24.7 37.6 ... 31.9 36.1 ... 44.0 35.6 19.7 40.2 ... ... ... 36.05405.79 Fe i 0.99 -1.85 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...5415.19 Fe i 4.39 0.51 188.5 163.3 179.9 185.6 175.6 172.9 163.6 144.5 153.8 173.8 172.0 164.7 179.3 133.7 189.9 170.15417.04 Fe i 4.42 -1.42 47.8 44.0 45.5 36.5 28.8 40.2 49.6 32.7 49.8 44.8 34.4 46.3 67.8 27.4 58.3 42.65434.53 Fe i 1.01 -2.12 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...5436.30 Fe i 4.39 -1.35 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...5464.29 Fe i 4.14 -1.62 78.2 81.2 59.0 ... 56.1 88.6 89.9 62.3 61.6 54.0 81.9 85.5 55.0 61.4 52.1 ...5470.09 Fe i 4.45 -1.60 24.2 28.7 ... 51.0 29.4 55.1 32.4 20.7 39.1 16.2 ... 30.9 22.1 ... 26.7 57.75501.48 Fe i 0.96 -3.05 250.9 257.7 261.3 265.2 254.2 274.4 297.6 242.7 230.1 273.8 279.0 263.5 260.7 238.4 274.6 ...5506.79 Fe i 0.99 -2.79 ... ... ... ... ... ... ... ... 275.1 ... ... ... 297.8 282.4 ... ...5539.29 Fe i 3.64 -2.59 75.1 72.4 73.9 82.0 61.6 74.2 95.8 58.1 60.1 54.3 78.0 71.8 56.2 36.9 78.9 86.25586.77 Fe i 3.37 -0.10 239.6 237.1 247.2 236.5 236.5 222.0 249.5 214.1 192.1 235.0 242.8 238.3 231.8 192.0 279.5 288.06120.26 Fe i 0.91 -5.94 78.6 78.7 89.7 107.2 105.7 96.8 113.2 88.6 75.8 100.8 97.2 82.1 82.5 63.2 87.8 112.16136.62 Fe i 2.45 -1.50 ... ... ... 288.4 ... ... ... ... ... ... ... ... ... 255.6 ... ...6137.00 Fe i 2.20 -2.95 ... ... ... 164.4 ... ... 174.1 ... ... ... ... ... ... 164.9 ... ...6151.62 Fe i 2.18 -3.37 134.8 134.6 126.9 147.0 136.5 142.0 141.4 124.7 137.7 139.2 133.4 150.7 125.6 115.7 147.1 146.26157.75 Fe i 4.07 -1.26 122.8 123.1 115.8 127.0 118.0 124.3 132.6 108.9 116.6 102.5 125.4 134.5 109.0 103.9 121.7 109.46159.38 Fe i 4.61 -1.97 16.6 ... 16.8 34.0 ... ... 40.8 19.3 ... 26.7 ... 16.9 ... ... ... ...6165.36 Fe i 4.14 -1.47 65.1 72.9 73.8 76.1 68.3 70.3 74.1 68.4 64.8 78.8 77.9 80.4 58.5 52.3 68.7 97.46173.34 Fe i 2.22 -2.85 167.9 171.2 159.8 156.2 160.9 166.2 179.1 166.7 159.6 165.3 179.0 179.5 150.8 156.8 177.6 175.46180.20 Fe i 2.73 -2.78 142.3 135.5 158.9 133.2 137.1 131.0 141.5 122.1 107.0 133.3 143.2 139.8 138.4 87.1 141.0 144.56187.99 Fe i 3.94 -1.58 80.3 86.8 97.2 85.8 78.7 80.5 85.4 77.2 66.0 78.5 77.8 70.5 78.7 60.8 92.8 94.06200.31 Fe i 2.61 -2.44 157.3 150.8 157.3 150.1 148.1 152.8 147.3 154.7 129.6 156.9 159.5 155.7 159.0 141.7 175.0 138.36213.43 Fe i 2.22 -2.66 179.4 184.1 192.5 182.8 191.9 187.1 183.4 196.9 200.1 189.3 198.5 187.0 168.8 177.8 200.1 184.16219.29 Fe i 2.20 -2.44 201.7 204.1 198.7 203.1 192.4 204.3 190.6 194.0 198.8 208.5 200.7 198.6 195.0 191.2 223.6 198.76226.74 Fe i 3.88 -2.20 71.7 72.4 49.0 70.7 75.1 69.2 70.5 49.4 53.9 59.6 76.6 69.8 53.1 47.5 51.8 87.66252.57 Fe i 2.40 -1.76 202.2 196.9 204.6 200.4 200.6 217.1 189.9 193.4 207.3 206.2 203.2 200.5 208.4 189.7 205.0 196.36265.13 Fe i 2.18 -2.55 181.0 179.9 190.6 183.7 167.1 189.0 201.1 177.6 181.4 192.7 182.8 173.4 175.5 171.9 188.3 213.46271.28 Fe i 3.32 -2.96 71.3 59.8 84.7 81.2 68.9 73.8 71.0 ... 56.2 58.1 55.0 63.8 60.7 50.9 51.9 65.46297.80 Fe i 2.22 -2.74 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6301.50 Fe i 3.65 -0.72 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6307.85 Fe i 3.64 -3.27 ... ... 26.1 26.6 ... ... ... ... ... 18.2 ... 21.8 ... ... ... ...6322.69 Fe i 2.59 -2.43 152.4 155.6 162.8 151.3 156.4 173.8 168.2 154.4 153.9 179.7 156.8 162.6 153.0 142.6 172.5 170.56330.85 Fe i 4.73 -1.22 51.9 63.2 45.0 54.3 37.3 49.6 43.7 35.7 34.3 33.5 52.2 ... 39.5 29.9 52.5 26.06335.33 Fe i 2.20 -2.23 196.9 216.0 218.7 213.1 216.5 233.3 228.9 206.0 211.5 208.0 226.1 201.3 203.5 202.1 209.1 233.06336.82 Fe i 3.69 -1.05 161.6 143.0 160.9 147.3 143.2 155.8 149.3 141.8 143.2 149.9 155.6 142.5 142.7 128.4 158.8 145.16344.15 Fe i 2.43 -2.92 187.9 191.9 179.6 195.9 185.9 191.7 200.8 170.1 170.1 174.4 189.4 196.4 159.6 145.3 191.2 204.86355.04 Fe i 2.84 -2.29 162.9 164.4 151.7 171.9 170.8 178.9 171.9 152.1 150.7 163.7 139.8 169.0 146.2 136.6 158.0 168.86380.75 Fe i 4.19 -1.50 97.3 102.2 89.6 77.1 87.3 90.2 103.2 85.2 83.9 89.3 75.3 93.9 77.4 68.0 87.3 85.86392.54 Fe i 2.28 -3.95 103.5 92.9 96.2 94.8 95.1 106.7 110.1 85.6 83.0 93.5 96.7 107.3 80.2 75.7 102.6 97.76393.61 Fe i 2.43 -1.63 228.9 245.7 256.1 259.2 238.7 263.7 256.5 239.2 233.3 253.0 248.8 238.4 236.2 231.7 264.9 280.46408.03 Fe i 3.69 -1.00 159.6 144.2 161.6 160.2 152.4 ... 173.0 129.1 125.9 160.8 158.2 156.8 146.4 140.1 167.5 175.06419.96 Fe i 4.73 -0.24 118.5 116.0 130.2 122.3 102.3 ... 118.8 109.5 111.5 92.4 135.7 117.2 116.6 108.2 120.5 124.76421.36 Fe i 2.28 -2.01 224.1 208.2 230.2 230.3 226.3 ... 228.7 222.5 209.4 229.7 235.3 226.7 221.9 212.6 233.4 272.46430.86 Fe i 2.18 -1.95 236.7 247.0 276.1 274.0 247.2 ... 293.5 245.6 256.8 260.1 252.5 264.6 248.3 228.3 235.1 ...6481.87 Fe i 2.27 -2.98 163.5 167.5 181.6 189.0 155.5 207.5 180.3 173.3 157.6 172.0 159.7 174.4 171.4 161.4 164.2 180.86498.94 Fe i 0.96 -4.69 177.3 190.6 201.9 216.9 187.1 219.8 220.2 183.3 181.2 197.4 192.1 191.6 163.4 173.1 183.1 235.56533.93 Fe i 4.55 -1.46 46.9 54.7 54.1 42.5 43.7 ... 47.6 46.8 48.5 54.7 61.4 51.9 50.7 41.8 49.1 37.16556.81 Fe i 4.79 -1.72 ... 16.9 23.8 24.3 ... 25.3 ... 22.8 25.6 19.6 ... 16.5 ... ... ... ...6569.22 Fe i 4.73 -0.42 126.2 122.2 124.0 117.1 109.5 139.8 124.8 98.3 111.5 96.3 128.5 130.1 100.2 88.2 101.6 128.66574.23 Fe i 0.99 -5.02 161.1 184.0 205.3 194.7 168.4 ... 200.7 159.5 153.3 190.8 196.9 180.7 171.2 ... ... ...6593.88 Fe i 2.43 -2.39 206.4 203.9 248.3 195.4 197.4 ... 238.5 162.8 128.0 215.0 220.9 202.0 196.8 171.6 231.3 249.66597.56 Fe i 4.79 -1.07 74.0 75.2 75.1 68.6 57.8 ... 73.9 55.4 35.0 66.5 72.9 74.8 67.0 43.8 69.7 78.26608.03 Fe i 2.28 -3.94 98.4 104.0 97.6 114.9 88.4 133.3 112.3 103.0 92.2 103.9 98.3 106.2 105.2 76.0 94.4 115.66609.12 Fe i 2.56 -2.66 175.7 175.4 164.2 169.7 167.0 198.7 177.0 166.1 195.2 172.0 156.1 161.2 162.4 165.2 158.6 174.16627.54 Fe i 4.54 -1.68 26.7 37.4 33.3 35.7 30.9 48.2 30.0 34.7 37.8 22.6 36.9 36.5 15.9 20.4 27.0 23.86633.76 Fe i 4.56 -0.82 125.1 120.0 126.9 118.1 113.7 148.6 115.4 110.8 118.0 108.7 134.8 122.3 97.5 93.4 126.1 105.06646.93 Fe i 2.60 -3.99 54.5 61.8 61.7 64.9 48.6 80.3 71.3 55.6 69.8 54.7 70.4 70.2 45.7 42.9 63.6 53.76653.85 Fe i 4.15 -2.52 23.7 22.9 39.1 35.4 26.0 41.2 44.8 34.8 41.6 ... 35.6 30.7 39.7 35.2 27.5 58.55414.08 Fe ii 3.22 -3.61 34.7 ... 39.0 ... 22.5 ... ... 25.8 30.9 ... 46.9 ... 44.2 ... 44.3 ...5425.25 Fe ii 3.20 -3.36 64.5 47.6 46.0 19.3 40.5 27.1 36.6 54.5 40.7 32.9 31.1 42.0 36.1 56.6 69.4 37.96149.25 Fe ii 3.89 -2.72 35.3 ... 30.5 29.9 36.7 30.4 ... 19.7 ... 29.8 ... 33.1 ... ... 31.3 38.36432.68 Fe ii 2.89 -3.71 62.8 48.8 58.5 37.8 51.7 ... 36.2 50.3 48.3 47.9 52.3 59.5 53.8 56.8 39.8 39.56456.39 Fe ii 3.90 -2.08 86.2 63.5 124.5 45.8 63.8 ... 78.3 22.8 ... 84.1 70.7 80.2 76.2 39.9 76.7 127.76320.43 La ii 0.17 -1.56 80.1 94.6 90.4 96.5 81.2 107.2 106.3 83.2 65.3 51.3 102.2 87.0 41.0 47.9 88.7 92.96390.46 La ii 0.32 -1.40 70.7 77.9 88.4 85.7 79.7 99.8 93.8 98.9 69.8 56.5 87.1 67.6 49.2 66.5 78.1 96.85528.41 Mg i 4.35 -0.36 217.6 216.7 209.4 211.1 200.0 231.6 234.9 215.0 205.8 215.0 234.1 214.4 204.7 209.8 220.6 230.1



λ(Å) elem χ gf 205 208 210 211 213 216 218 221 227 228 229 233 239 242 247 2506318.72 Mg i 5.11 -1.97 46.5 41.8 55.4 46.9 39.1 51.1 47.3 ... ... 31.8 35.8 44.6 ... 41.6 ... 58.66319.24 Mg i 5.11 -2.21 ... 36.2 ... 23.7 ... 25.3 33.9 19.2 29.8 ... ... ... 19.4 ... 43.4 34.66319.49 Mg i 5.11 -2.43 ... ... 35.7 ... ... ... ... ... ... ... ... 24.0 ... 36.0 ... ...5420.36 Mn i 2.14 -1.46 222.1 201.1 213.4 230.0 191.8 206.6 221.1 185.9 155.9 217.5 222.2 204.7 224.6 180.7 226.2 211.85432.55 Mn i 0.00 -3.80 248.6 272.1 276.2 294.5 270.2 297.9 ... 237.3 266.0 278.9 274.8 271.8 257.9 239.2 281.5 ...5516.77 Mn i 2.18 -1.85 158.4 ... 175.1 172.6 141.3 191.9 182.3 140.3 160.1 ... 178.1 160.3 126.4 131.3 ... 157.86154.23 Na i 2.10 -1.56 37.4 38.0 36.5 40.5 ... 42.5 27.2 ... 33.8 16.6 29.5 37.6 16.8 ... 24.6 44.56160.75 Na i 2.10 -1.26 50.6 51.6 51.1 57.3 40.1 65.8 60.3 27.5 48.6 29.5 51.5 63.1 ... 28.9 54.0 88.75416.38 Nd ii 0.86 -0.98 ... 38.9 50.4 43.3 36.4 40.5 40.0 ... 34.3 ... 21.4 22.1 22.8 ... 57.5 65.45431.54 Nd ii 1.12 -0.47 65.2 60.6 57.3 68.0 ... 66.6 73.5 69.4 90.6 47.5 64.9 52.6 30.8 66.1 86.5 66.95485.71 Nd ii 1.26 -0.12 37.2 40.6 45.4 54.2 41.9 58.1 54.4 50.7 36.4 ... 53.4 51.7 25.4 34.8 ... 36.75578.73 Ni i 1.68 -2.67 142.9 143.8 160.6 142.6 142.5 148.6 164.2 121.6 123.3 150.8 145.0 142.7 129.9 129.3 135.9 184.65587.87 Ni i 1.93 -2.37 171.5 168.7 191.9 154.0 154.7 177.2 172.2 131.3 166.7 165.8 184.2 168.3 163.9 121.0 180.6 177.05589.37 Ni i 3.90 -1.15 27.6 35.4 31.8 32.0 26.8 28.3 37.3 ... 26.3 20.2 20.1 26.7 19.7 17.8 ... ...5593.75 Ni i 3.90 -0.79 51.5 56.2 59.7 43.0 35.4 39.2 43.3 32.4 45.5 35.0 38.1 51.7 29.4 38.9 45.0 49.46128.97 Ni i 1.68 -3.39 78.6 85.2 93.6 97.5 93.8 106.0 104.8 90.8 71.8 92.9 90.9 99.0 78.8 79.3 90.2 93.86130.14 Ni i 4.27 -0.98 17.0 19.7 ... 19.2 ... ... 33.2 ... ... ... ... ... ... ... 25.6 ...6177.25 Ni i 1.83 -3.60 57.7 68.3 77.5 73.0 58.1 79.0 61.1 54.6 55.8 61.7 76.6 66.2 60.1 45.9 78.2 63.96186.72 Ni i 4.11 -0.90 33.5 33.6 40.0 39.3 38.9 50.3 52.2 29.5 22.0 30.9 32.3 30.2 25.1 25.5 41.4 86.06204.61 Ni i 4.09 -1.15 28.4 23.7 ... 20.2 25.2 ... 22.5 16.8 23.9 ... 25.1 27.2 27.4 ... 30.5 ...6223.99 Ni i 4.10 -0.97 45.0 44.0 41.2 44.4 34.7 28.8 44.0 39.0 ... 25.2 40.3 32.3 26.2 31.1 48.3 61.96230.10 Ni i 4.11 -1.20 26.8 30.5 39.3 ... 26.7 45.1 ... 15.2 ... 28.3 ... 32.0 ... 24.3 27.2 46.66322.17 Ni i 4.15 -1.21 23.4 26.4 16.4 29.0 21.4 ... ... ... ... ... ... ... 15.2 ... ... ...6327.60 Ni i 1.68 -3.09 121.5 118.8 124.0 125.6 103.0 131.3 129.8 116.6 109.6 112.7 116.5 109.8 110.1 97.3 129.5 130.56378.26 Ni i 4.15 -0.82 33.2 25.4 23.3 40.5 33.4 32.9 50.0 30.1 31.3 30.8 32.5 34.2 22.8 ... 38.5 47.26384.67 Ni i 4.15 -1.00 ... ... 29.8 47.9 ... 38.7 45.6 ... 38.8 ... 37.7 47.0 ... 33.3 58.1 43.26482.80 Ni i 1.94 -2.85 103.8 99.7 121.0 116.2 86.2 127.1 117.2 104.6 101.6 105.6 97.7 113.8 103.5 100.4 110.0 125.86586.32 Ni i 1.95 -2.79 142.5 165.3 179.6 139.6 154.3 ... 157.2 80.0 58.3 150.4 174.5 155.2 154.8 120.3 185.7 185.86598.61 Ni i 4.24 -0.93 ... 30.9 ... ... ... ... 33.8 17.0 ... 17.8 ... ... ... ... 42.1 ...6635.14 Ni i 4.42 -0.75 ... 17.0 ... ... ... 36.7 ... 31.8 23.1 ... ... ... ... ... ... 20.16300.31 O i 0.00 -9.75 ... ... ... ... ... ... ... ... ... ... ... ... ... 64.4 ... 65.46363.79 O i 0.02 -10.25 ... 39.7 39.4 47.0 44.8 51.8 51.7 34.5 39.3 42.8 50.7 31.8 27.7 37.2 41.5 54.85526.82 Sc ii 1.77 0.03 112.9 110.3 103.8 104.6 115.0 123.4 99.0 108.6 113.0 100.1 117.8 102.3 118.4 101.7 114.5 103.86245.62 Sc ii 1.51 -0.97 80.1 72.5 64.2 69.6 77.7 ... 69.7 67.1 76.2 64.0 67.8 65.6 61.8 65.4 74.5 69.56309.90 Sc ii 1.50 -1.52 33.8 ... ... ... 46.4 58.7 57.0 ... 57.1 22.3 44.7 47.9 ... ... ... ...6604.60 Sc ii 1.36 -1.31 87.2 98.9 97.4 100.1 93.4 111.0 98.2 109.9 84.6 92.1 85.0 89.8 87.5 78.4 75.7 95.76125.03 Si i 5.62 -1.57 41.7 ... ... 48.4 ... ... ... ... 16.4 ... ... ... ... ... ... ...6142.48 Si i 5.62 -1.51 ... ... ... ... ... ... ... ... 18.5 ... 21.8 ... ... ... ... ...6145.02 Si i 5.61 -1.37 ... ... ... ... 27.6 19.2 ... ... ... 23.1 34.5 ... ... ... ... ...6155.14 Si i 5.62 -0.80 57.4 45.8 40.5 45.3 43.8 51.5 50.1 43.2 54.9 31.3 42.5 45.8 32.7 45.0 42.3 33.16237.33 Si i 5.61 -1.02 32.5 28.2 18.4 23.9 21.0 24.3 18.1 23.3 25.5 ... 27.3 19.1 23.2 23.2 ... 37.46243.82 Si i 5.61 -1.27 41.7 47.4 35.3 56.3 29.1 ... ... ... 29.2 26.3 46.1 33.8 23.5 ... 37.5 40.25490.16 Ti i 1.46 -0.93 100.1 107.3 122.6 134.3 112.7 135.9 148.3 109.9 112.4 105.2 123.3 134.1 93.7 95.3 133.0 140.35503.90 Ti i 2.58 -0.19 67.3 70.1 71.8 83.0 68.6 88.7 89.3 57.4 44.5 58.7 83.7 79.9 43.2 39.5 71.4 109.56126.22 Ti i 1.07 -1.42 122.8 150.5 134.6 164.4 145.3 169.4 170.1 134.9 130.7 143.2 151.8 140.1 118.9 117.5 159.2 189.26220.50 Ti i 2.68 -0.14 ... ... ... 87.8 ... ... 86.3 ... ... 52.1 ... ... ... ... ... 65.26258.10 Ti i 1.44 -0.35 151.1 182.8 187.5 207.9 184.1 228.0 232.0 170.1 199.9 212.6 206.9 178.7 150.3 142.5 153.9 297.96303.77 Ti i 1.44 -1.57 77.5 89.4 88.6 118.9 93.1 118.0 115.2 76.7 78.8 86.2 112.3 91.1 80.3 67.8 91.0 122.56312.24 Ti i 1.46 -1.55 ... 87.6 93.7 109.2 ... ... ... 73.4 76.1 92.9 ... 91.8 61.6 66.7 74.4 130.06336.10 Ti i 1.44 -1.74 67.5 80.4 76.9 94.4 80.4 96.4 98.4 69.0 79.1 77.2 84.1 84.7 55.2 40.0 84.6 128.36508.12 Ti i 1.43 -2.05 31.5 45.7 59.1 81.0 57.5 86.4 84.7 45.7 66.1 52.5 66.3 65.7 26.5 ... 46.9 110.66556.08 Ti i 1.46 -1.07 ... ... ... ... 133.0 ... ... ... ... ... ... ... ... ... ... ...6599.13 Ti i 0.90 -2.09 111.3 120.8 153.8 181.9 142.2 ... 186.1 101.0 113.8 148.0 156.7 132.8 110.5 99.0 149.9 221.26666.53 Ti i 1.46 -1.62 19.2 34.5 34.7 36.1 ... 41.7 38.9 ... ... ... 35.8 18.4 15.1 ... 37.9 59.05418.77 Ti ii 1.58 -2.11 111.5 96.6 114.0 97.2 111.6 101.5 97.4 103.3 95.0 100.9 109.3 104.3 82.5 80.7 98.5 99.36219.94 Ti ii 2.06 -2.82 36.1 ... ... ... ... ... ... ... ... ... 41.7 37.2 ... 28.1 40.3 ...6559.58 Ti ii 2.05 -2.02 ... ... ... ... 49.2 72.9 54.2 ... ... ... ... ... ... ... ... ...6606.95 Ti ii 2.06 -2.79 45.7 49.4 43.0 52.7 40.4 67.7 35.3 55.5 45.0 40.3 25.8 51.0 34.4 48.0 27.7 23.86680.13 Ti ii 3.09 -1.86 36.3 ... 29.2 19.9 22.4 25.4 18.2 26.7 ... 28.0 21.7 27.1 ... ... 31.5 ...6119.53 V i 1.06 -0.32 ... ... 96.7 ... ... ... ... ... ... ... ... ... ... ... ... ...6128.33 V i 1.05 -2.30 ... ... ... ... ... ... ... ... ... ... ... 17.8 ... ... ... 41.66135.37 V i 1.05 -0.75 80.7 100.2 104.6 133.1 102.6 136.0 140.1 89.2 106.0 114.3 122.1 104.1 ... 87.5 89.4 150.56150.15 V i 0.30 -1.79 120.7 126.9 138.9 158.2 148.7 173.3 182.5 122.2 151.2 162.3 152.9 149.4 112.1 106.3 147.3 198.56199.19 V i 0.29 -1.29 155.0 166.1 195.5 221.2 177.2 217.9 212.1 154.2 177.1 186.7 215.9 196.5 144.0 141.5 186.4 232.66216.36 V i 0.28 -0.81 150.6 174.7 188.8 200.7 183.0 219.5 213.9 166.7 204.3 195.9 198.9 195.9 152.6 162.2 180.1 225.76224.51 V i 0.29 -2.01 107.7 122.9 132.0 159.3 142.3 154.2 161.2 120.2 ... 130.9 148.4 140.2 91.9 111.0 141.8 196.76233.20 V i 0.28 -2.07 66.4 91.8 95.3 108.2 88.2 128.0 134.0 94.6 99.7 101.9 112.8 94.1 63.2 69.7 81.0 130.66243.11 V i 0.30 -0.98 ... ... ... ... ... ... ... ... ... ... ... 281.0 ... ... ... ...6251.82 V i 0.29 -1.30 112.2 139.4 140.5 169.3 140.3 177.0 160.9 137.1 155.4 156.3 161.9 141.1 116.6 117.8 135.5 164.76274.66 V i 0.27 -1.67 ... 111.8 126.9 142.9 ... ... ... 109.8 118.2 130.0 ... ... 92.2 79.0 107.5 156.56357.29 V i 1.85 -0.91 17.3 27.6 33.7 51.5 27.1 45.6 49.2 15.4 16.9 26.8 36.8 39.5 ... 20.6 34.0 69.36452.32 V i 1.19 -1.21 68.2 80.4 114.3 123.4 90.4 ... 137.6 62.4 64.5 81.1 103.8 98.1 69.6 45.7 ... 167.86504.19 V i 1.18 -1.23 38.1 42.8 50.1 73.6 60.8 87.5 86.4 49.6 50.1 44.5 72.3 50.7 45.2 41.3 39.0 73.26531.41 V i 1.22 -0.84 ... ... ... ... 70.3 ... 97.2 ... ... ... 84.6 ... ... ... ... ...5402.78 Y ii 1.84 -0.51 44.6 24.4 45.0 58.4 53.7 55.9 55.2 40.5 50.7 29.3 26.5 60.0 ... 32.0 74.5 49.26362.35 Zn i 5.80 0.14 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6127.48 Zr i 0.15 -1.06 63.0 81.6 92.7 122.1 84.2 114.8 132.5 93.4 78.8 68.7 104.0 87.2 45.2 56.2 80.3 132.56140.46 Zr i 0.52 -1.41 ... 16.4 ... 27.9 ... 31.8 35.0 15.9 ... ... 28.7 ... ... ... ... 50.96143.18 Zr i 0.07 -1.10 59.9 77.9 75.0 121.3 79.9 127.9 113.2 83.8 98.1 84.3 110.9 85.8 52.4 59.6 67.4 136.36192.95 Zr i 0.54 -2.07 ... 17.4 20.1 21.9 15.5 21.4 19.6 ... ... ... ... ... ... ... ... 34.6



Equivalent witdh, one star per column, BLxxxλ(Å) elem χ gf 253 257 258 260 261 262 266 267 269 278 279 295 300 304 311 315 3236141.73 Ba ii 0.70 -0.08 287.8 ... 297.3 232.5 218.4 230.4 173.6 248.7 246.2 ... 167.2 ... 293.7 261.3 253.0 215.4 260.36496.91 Ba ii 0.60 -0.38 282.5 290.1 ... 228.3 202.2 266.1 185.8 237.1 273.8 ... 160.6 ... 265.7 244.0 ... 193.9 264.96122.23 Ca i 1.89 -0.32 279.3 291.0 281.2 242.9 236.3 ... 182.2 223.9 258.7 ... 174.8 ... 275.1 271.8 247.5 201.6 252.66156.03 Ca i 2.52 -2.39 38.4 46.8 43.7 32.1 20.5 ... ... 32.0 28.5 45.4 ... ... 17.5 15.6 36.5 ... ...6161.30 Ca i 2.52 -1.27 163.2 166.7 180.6 127.1 127.9 113.3 50.6 132.8 132.2 177.4 47.0 161.4 145.4 143.2 138.3 87.7 148.16162.17 Ca i 1.90 -0.32 ... ... ... 278.8 266.5 256.1 196.0 270.4 282.6 ... 184.9 ... 291.6 288.4 277.4 221.8 282.16166.44 Ca i 2.52 -1.14 148.6 150.4 150.5 136.3 125.4 121.0 64.0 146.7 124.0 147.9 37.8 157.0 135.9 130.7 128.6 103.2 113.86169.04 Ca i 2.52 -0.80 163.3 169.7 161.6 145.9 135.6 129.1 94.4 148.5 144.6 159.9 64.5 154.9 152.0 149.4 133.3 123.6 105.86169.56 Ca i 2.52 -0.48 170.4 196.1 182.1 165.5 140.1 154.1 112.4 168.4 151.7 177.3 86.8 196.2 176.1 171.7 171.3 156.6 146.56439.08 Ca i 2.52 0.39 236.5 223.4 247.4 213.6 209.2 230.1 183.4 225.8 215.5 238.4 155.1 260.0 224.0 228.4 ... 207.2 235.26455.60 Ca i 2.52 -1.29 84.2 113.4 187.4 152.1 ... ... 29.3 153.2 75.5 170.9 ... 192.0 ... 130.8 ... 155.9 60.66471.67 Ca i 2.52 -0.76 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6493.79 Ca i 2.52 -0.32 209.7 223.2 189.5 187.3 182.2 181.5 139.8 196.4 194.6 211.4 126.6 199.2 210.3 205.2 204.5 156.9 195.86499.65 Ca i 2.52 -0.82 144.6 159.2 138.6 118.5 121.0 133.3 77.6 106.5 138.1 151.2 45.9 141.8 127.3 131.3 152.9 105.2 126.86508.84 Ca i 2.52 -2.41 45.2 57.3 55.8 38.7 32.4 43.9 ... 37.1 46.5 61.5 ... 57.0 29.9 28.7 38.6 ... 41.56330.09 Cr i 0.94 -2.92 152.8 162.8 163.2 142.5 123.5 130.7 64.1 132.6 130.4 172.1 35.2 159.7 150.2 140.8 137.0 106.9 158.76645.13 Eu ii 1.37 0.20 75.8 74.3 90.6 51.4 62.8 54.1 30.6 64.9 82.4 78.2 34.6 82.9 61.2 59.4 113.1 75.5 70.15369.96 Fe i 4.37 0.54 171.7 171.0 180.1 182.4 172.0 179.4 112.4 164.9 141.9 247.5 101.7 208.0 164.4 158.8 238.5 145.2 165.15383.37 Fe i 4.31 0.50 153.4 169.7 181.0 171.1 145.9 149.5 130.6 135.1 153.1 182.5 93.5 173.8 187.3 160.0 139.8 129.9 162.65386.34 Fe i 4.16 -1.74 ... 65.9 ... ... 45.7 33.1 ... 56.6 55.8 ... 17.7 57.3 ... 35.6 39.6 58.9 35.85393.17 Fe i 3.24 -0.92 226.5 248.5 235.0 210.7 213.9 192.2 169.3 229.2 204.2 235.2 125.3 246.4 218.6 224.3 198.2 210.6 178.55395.22 Fe i 4.45 -1.73 37.3 39.0 43.7 ... ... 17.1 ... 31.1 27.5 31.5 ... 36.2 ... 31.6 29.1 31.2 22.25405.79 Fe i 0.99 -1.85 ... ... ... ... ... ... 283.9 ... ... ... 284.6 ... ... ... ... ... ...5415.19 Fe i 4.39 0.51 177.9 181.1 190.1 182.6 159.5 166.7 145.1 161.7 146.3 173.3 110.4 178.1 163.2 166.8 141.8 150.5 151.35417.04 Fe i 4.42 -1.42 65.8 47.7 43.5 36.1 39.1 43.0 19.2 77.1 23.1 33.3 ... 46.4 35.3 32.4 ... 40.8 23.05434.53 Fe i 1.01 -2.12 ... ... ... ... ... ... 271.4 ... ... ... 259.8 ... ... ... ... ... ...5436.30 Fe i 4.39 -1.35 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...5464.29 Fe i 4.14 -1.62 92.2 ... 90.0 56.4 59.3 56.0 ... ... 74.2 ... ... 108.1 71.5 58.9 62.9 19.7 67.05470.09 Fe i 4.45 -1.60 34.5 46.1 ... 28.5 33.1 23.5 ... 29.9 31.6 43.6 ... 35.5 33.3 29.8 40.3 ... 39.25501.48 Fe i 0.96 -3.05 277.6 268.8 274.5 250.1 264.8 245.2 207.2 250.1 234.0 279.6 185.4 278.9 277.9 276.9 241.0 237.5 285.35506.79 Fe i 0.99 -2.79 ... ... ... ... ... 290.9 212.0 ... ... ... 204.9 ... ... ... 294.6 271.7 ...5539.29 Fe i 3.64 -2.59 78.5 95.8 83.4 71.9 61.5 55.6 26.6 87.2 52.8 86.6 21.8 84.0 ... 57.8 48.2 48.9 61.55586.77 Fe i 3.37 -0.10 252.5 248.6 278.2 246.4 225.6 213.5 171.9 224.0 199.3 243.8 149.2 274.9 225.7 221.1 190.1 212.5 211.96120.26 Fe i 0.91 -5.94 101.4 111.6 106.3 102.5 97.4 72.3 28.4 93.0 83.3 96.8 36.1 102.3 94.0 107.8 106.2 76.8 82.36136.62 Fe i 2.45 -1.50 ... ... ... ... 254.6 ... 237.6 274.3 282.0 ... 159.6 ... ... ... ... 222.3 ...6137.00 Fe i 2.20 -2.95 ... ... ... ... 172.1 ... ... 158.3 163.6 195.8 85.9 ... ... ... ... 139.1 ...6151.62 Fe i 2.18 -3.37 148.8 136.1 144.6 122.5 127.2 128.9 90.8 124.8 122.8 137.6 76.5 149.9 124.0 140.5 128.3 119.4 134.36157.75 Fe i 4.07 -1.26 129.7 142.4 139.5 102.0 100.3 125.4 59.8 125.9 107.2 143.3 43.0 146.4 127.3 109.5 119.9 107.2 130.06159.38 Fe i 4.61 -1.97 27.0 41.7 37.7 20.1 ... 29.4 ... 30.5 25.2 52.9 ... 45.3 ... ... 22.7 ... ...6165.36 Fe i 4.14 -1.47 82.7 87.5 86.5 59.3 65.0 67.8 39.1 75.4 60.6 84.0 28.0 76.3 65.3 68.5 72.3 80.4 72.86173.34 Fe i 2.22 -2.85 183.2 186.7 180.4 163.0 161.3 156.7 100.5 170.7 148.2 171.0 99.7 184.4 179.6 176.4 169.7 143.7 190.26180.20 Fe i 2.73 -2.78 ... 148.7 168.4 147.2 165.3 100.1 98.9 146.3 125.3 134.8 84.8 157.8 140.4 138.1 131.9 125.8 76.96187.99 Fe i 3.94 -1.58 87.3 108.2 97.1 96.2 67.7 81.1 42.9 85.6 75.3 94.7 36.8 78.5 78.1 88.2 80.3 78.4 77.86200.31 Fe i 2.61 -2.44 171.8 164.5 170.2 165.1 151.3 147.9 114.6 166.8 145.1 149.0 104.1 165.9 163.2 161.1 157.6 139.1 146.96213.43 Fe i 2.22 -2.66 196.8 195.7 192.3 193.3 161.9 200.9 140.1 184.5 185.9 204.3 134.1 197.2 188.4 201.4 183.5 151.7 206.26219.29 Fe i 2.20 -2.44 215.8 211.3 206.0 214.4 173.8 206.8 154.0 197.0 189.1 201.4 132.3 200.5 189.8 202.1 192.9 179.9 213.16226.74 Fe i 3.88 -2.20 67.3 78.6 72.2 63.4 ... 61.1 22.7 63.5 46.8 ... ... 80.6 52.6 60.6 54.8 57.7 69.86252.57 Fe i 2.40 -1.76 217.6 211.9 205.2 197.7 197.9 202.1 168.2 189.9 202.8 213.3 135.4 221.5 196.8 218.8 205.8 173.9 217.16265.13 Fe i 2.18 -2.55 195.3 207.8 209.4 188.2 163.6 180.2 107.0 180.3 187.9 208.3 120.1 206.1 191.0 195.1 193.5 160.3 175.76271.28 Fe i 3.32 -2.96 66.5 82.4 59.1 64.5 52.9 55.9 29.5 59.5 57.6 70.0 37.0 75.4 44.6 57.2 65.8 47.7 51.96297.80 Fe i 2.22 -2.74 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6301.50 Fe i 3.65 -0.72 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6307.85 Fe i 3.64 -3.27 ... 25.0 ... ... ... 19.5 ... 21.4 ... 27.9 ... ... ... ... ... ... ...6322.69 Fe i 2.59 -2.43 176.2 160.4 163.1 176.9 148.3 153.4 124.0 153.4 155.5 167.8 99.6 169.0 144.1 158.5 165.5 142.7 166.66330.85 Fe i 4.73 -1.22 35.3 ... 50.8 42.8 ... 49.4 ... 41.6 36.5 54.1 ... 50.0 37.7 28.8 ... 27.6 40.76335.33 Fe i 2.20 -2.23 235.2 224.2 214.1 221.1 207.0 213.7 161.9 201.8 201.1 225.6 151.8 230.9 213.8 237.8 206.4 180.8 234.26336.82 Fe i 3.69 -1.05 157.5 150.6 146.2 153.3 144.4 134.0 122.5 162.7 140.0 156.0 80.9 161.6 152.4 155.2 146.1 92.0 148.46344.15 Fe i 2.43 -2.92 206.9 188.4 191.4 183.1 161.7 208.7 109.9 165.4 163.5 209.6 82.6 217.8 189.7 184.2 175.9 153.5 176.26355.04 Fe i 2.84 -2.29 189.5 177.7 174.5 181.2 136.5 149.6 96.0 159.7 152.6 186.8 100.9 183.4 144.1 166.6 154.4 144.1 162.06380.75 Fe i 4.19 -1.50 106.5 99.1 89.6 81.8 94.7 77.4 38.4 99.5 76.0 98.8 34.6 93.0 78.3 83.4 93.6 59.7 87.26392.54 Fe i 2.28 -3.95 93.9 98.2 104.3 108.8 100.9 87.8 30.4 99.3 85.8 101.4 36.9 113.5 87.1 106.7 96.7 67.5 86.16393.61 Fe i 2.43 -1.63 263.2 261.0 261.7 246.0 246.5 249.8 169.9 240.0 240.2 250.0 161.7 273.4 256.6 252.2 253.5 198.7 243.86408.03 Fe i 3.69 -1.00 147.4 168.0 168.1 153.0 161.0 150.2 102.9 144.7 123.7 175.4 77.8 173.2 153.4 158.4 ... 154.9 147.66419.96 Fe i 4.73 -0.24 125.2 148.8 137.2 124.9 99.9 116.3 59.9 131.9 112.6 133.6 51.8 127.7 98.9 126.0 ... 99.8 110.06421.36 Fe i 2.28 -2.01 251.2 235.6 232.1 235.2 221.2 229.9 178.8 205.2 225.5 236.9 152.3 248.4 222.6 236.1 ... 204.5 240.86430.86 Fe i 2.18 -1.95 278.9 292.2 274.9 270.0 255.0 274.3 195.8 250.8 263.4 294.1 165.1 269.1 240.9 279.6 ... 202.8 263.76481.87 Fe i 2.27 -2.98 184.2 195.7 161.7 172.9 160.0 178.4 142.7 163.6 172.2 197.0 97.8 179.3 175.8 163.6 197.4 151.7 181.36498.94 Fe i 0.96 -4.69 215.0 211.0 213.0 200.6 186.2 192.6 137.8 180.5 203.4 239.2 97.4 219.5 202.4 190.9 228.8 152.8 211.56533.93 Fe i 4.55 -1.46 55.0 45.7 61.3 52.6 52.5 55.8 ... 50.3 42.8 50.6 ... 57.1 45.8 41.9 ... 35.9 50.16556.81 Fe i 4.79 -1.72 24.7 26.9 28.1 ... 26.3 37.4 ... 19.1 31.8 23.4 ... ... 17.2 15.5 58.1 ... 20.56569.22 Fe i 4.73 -0.42 123.7 136.9 131.4 109.9 96.8 133.7 70.0 106.8 116.7 135.7 51.0 139.2 100.6 109.3 144.0 97.8 119.66574.23 Fe i 0.99 -5.02 188.7 190.3 208.8 185.8 165.0 ... 117.7 176.4 ... 189.1 83.8 ... 183.9 197.3 ... 143.8 193.66593.88 Fe i 2.43 -2.39 190.6 210.9 225.2 208.0 174.9 160.8 124.7 193.5 138.1 207.2 90.9 241.2 213.2 200.8 ... 172.9 196.06597.56 Fe i 4.79 -1.07 51.5 73.6 85.0 62.9 46.7 62.5 19.9 66.6 75.0 56.3 19.4 74.2 64.4 55.1 ... 66.1 55.26608.03 Fe i 2.28 -3.94 117.4 112.9 115.0 111.9 87.4 105.7 54.9 78.3 131.0 120.0 59.3 119.9 97.5 100.9 155.1 73.8 110.76609.12 Fe i 2.56 -2.66 157.1 176.1 163.4 166.0 128.7 166.6 120.0 152.2 196.8 176.8 76.7 166.4 158.0 162.5 223.2 122.1 190.26627.54 Fe i 4.54 -1.68 53.4 32.8 41.7 36.0 30.4 43.2 18.5 29.8 56.9 42.9 ... 25.8 33.0 20.6 105.3 24.3 29.66633.76 Fe i 4.56 -0.82 128.0 130.5 131.2 111.8 107.8 127.5 49.9 136.3 141.5 134.0 44.8 133.6 98.7 103.5 149.8 98.2 114.26646.93 Fe i 2.60 -3.99 78.9 74.4 86.7 54.1 65.0 60.6 17.9 66.5 53.2 76.2 ... 79.9 69.6 71.1 86.1 52.4 51.26653.85 Fe i 4.15 -2.52 ... 44.4 46.5 49.4 29.9 38.5 ... 31.8 42.7 40.8 16.1 50.9 21.0 37.8 ... 29.9 34.45414.08 Fe ii 3.22 -3.61 ... ... ... ... 20.3 22.5 ... 30.9 ... ... ... ... ... ... ... ... ...5425.25 Fe ii 3.20 -3.36 35.6 ... 45.5 52.5 82.1 51.8 51.2 58.9 ... 30.6 56.7 30.4 40.7 29.2 20.7 70.4 22.66149.25 Fe ii 3.89 -2.72 35.2 28.9 35.5 35.8 ... 34.0 25.9 27.7 20.0 ... ... 43.4 28.2 34.0 38.6 24.7 ...6432.68 Fe ii 2.89 -3.71 49.8 26.0 46.2 65.1 55.5 74.1 56.4 50.3 62.1 41.8 33.7 ... 47.1 45.9 ... 40.4 47.16456.39 Fe ii 3.90 -2.08 40.4 31.9 64.7 ... 60.6 79.2 29.5 86.6 ... 71.4 49.2 86.6 71.8 67.2 ... 66.2 37.76320.43 La ii 0.17 -1.56 94.8 128.7 111.5 63.0 70.7 53.0 20.2 91.6 73.2 109.7 25.4 130.3 118.1 82.0 79.3 63.4 68.76390.46 La ii 0.32 -1.40 108.4 109.9 98.9 68.4 64.9 69.4 30.2 72.0 80.7 95.8 24.6 108.7 93.6 ... 74.5 79.4 70.45528.41 Mg i 4.35 -0.36 227.2 223.4 220.8 205.5 219.7 217.4 169.5 216.3 221.3 234.8 164.9 225.7 240.7 208.9 208.1 207.0 220.0



λ(Å) elem χ gf 253 257 258 260 261 262 266 267 269 278 279 295 300 304 311 315 3236318.72 Mg i 5.11 -1.97 44.0 64.6 47.2 32.4 25.6 50.3 27.4 43.6 28.3 ... ... 49.7 42.8 ... 41.7 ... 30.76319.24 Mg i 5.11 -2.21 26.0 25.8 ... ... ... 20.5 ... ... 26.4 38.8 ... 41.9 ... ... 25.7 ... ...6319.49 Mg i 5.11 -2.43 ... ... ... ... ... ... ... 28.2 ... ... ... ... 31.0 18.5 ... ... 15.25420.36 Mn i 2.14 -1.46 203.5 220.4 238.5 190.0 185.7 193.4 71.1 201.4 204.6 222.2 57.3 229.3 196.5 217.0 172.2 199.3 192.75432.55 Mn i 0.00 -3.80 292.1 ... 283.8 268.2 ... 271.9 156.3 236.9 254.0 ... 111.4 ... 281.0 270.9 284.1 220.5 291.95516.77 Mn i 2.18 -1.85 167.7 174.7 179.1 146.2 156.5 142.5 49.9 139.3 164.4 210.9 ... 211.2 158.8 161.1 154.3 108.0 144.26154.23 Na i 2.10 -1.56 41.9 37.6 45.3 ... 25.0 ... ... 31.7 24.5 52.8 16.2 48.8 33.4 17.0 24.3 ... ...6160.75 Na i 2.10 -1.26 68.6 62.1 63.6 34.0 32.5 25.2 ... 58.8 51.6 96.0 ... 70.2 52.3 38.9 47.2 26.3 44.35416.38 Nd ii 0.86 -0.98 19.8 36.1 54.3 40.6 48.5 ... ... 37.3 18.9 32.6 ... 69.6 35.0 33.4 ... ... ...5431.54 Nd ii 1.12 -0.47 70.6 75.9 121.6 50.1 40.4 59.3 ... 56.7 61.9 74.7 27.7 91.0 60.1 35.5 ... 38.3 58.75485.71 Nd ii 1.26 -0.12 55.3 60.8 59.8 31.8 46.5 37.0 ... 23.3 35.6 66.5 ... 50.6 59.8 43.1 34.8 ... 33.95578.73 Ni i 1.68 -2.67 147.4 165.4 161.2 125.8 121.5 121.0 96.1 148.4 141.2 174.5 81.1 161.9 143.0 152.3 131.7 128.5 140.85587.87 Ni i 1.93 -2.37 171.3 204.0 176.8 159.7 138.7 142.1 79.7 164.8 145.6 188.1 73.0 191.4 161.9 177.5 131.3 148.8 146.65589.37 Ni i 3.90 -1.15 37.4 25.3 25.2 38.7 28.7 22.6 ... ... 40.2 39.9 ... 36.7 ... 24.5 31.2 28.5 25.95593.75 Ni i 3.90 -0.79 39.7 53.6 47.5 38.2 45.3 42.2 ... 28.6 30.1 54.8 39.7 43.5 35.9 35.2 19.7 49.5 24.36128.97 Ni i 1.68 -3.39 105.2 96.0 99.1 93.3 86.7 94.7 48.8 92.7 86.6 108.3 34.0 99.2 99.3 95.3 77.0 42.2 93.86130.14 Ni i 4.27 -0.98 16.7 27.5 ... ... 25.4 ... ... 26.6 ... ... ... 22.4 ... ... ... 25.1 17.26177.25 Ni i 1.83 -3.60 82.2 80.4 76.0 62.4 69.3 61.0 39.8 69.0 56.1 66.7 31.7 78.2 72.8 67.1 67.9 66.2 62.46186.72 Ni i 4.11 -0.90 37.1 52.5 52.6 37.1 ... 20.1 ... 27.4 25.2 44.9 23.9 37.3 39.8 ... 33.6 32.3 30.76204.61 Ni i 4.09 -1.15 40.7 ... ... 15.4 ... 17.3 16.5 16.9 ... 26.0 15.8 20.8 ... ... 21.7 31.3 15.06223.99 Ni i 4.10 -0.97 37.8 45.9 49.3 23.0 39.9 ... ... ... 25.7 46.1 ... 31.2 28.4 30.1 42.3 42.0 ...6230.10 Ni i 4.11 -1.20 35.0 46.7 43.8 ... ... ... ... 36.6 29.1 70.0 ... 44.2 31.1 31.5 ... 21.9 47.86322.17 Ni i 4.15 -1.21 ... 32.3 ... ... ... ... 23.1 ... ... ... ... ... 20.4 24.4 ... ... ...6327.60 Ni i 1.68 -3.09 125.6 135.2 130.8 115.7 115.6 117.1 71.5 113.1 110.2 130.9 43.0 133.5 116.1 121.3 118.2 92.1 114.86378.26 Ni i 4.15 -0.82 40.4 46.7 35.4 30.4 18.4 32.6 21.6 23.9 29.3 46.1 ... 41.6 26.0 20.4 40.0 41.3 31.96384.67 Ni i 4.15 -1.00 38.5 49.3 60.2 ... ... 20.0 18.8 46.5 47.6 50.0 30.6 54.0 44.0 28.1 38.4 ... 34.26482.80 Ni i 1.94 -2.85 113.7 115.3 110.1 98.6 93.3 118.0 66.3 110.8 136.0 134.6 57.3 110.1 102.2 96.2 136.9 93.3 105.56586.32 Ni i 1.95 -2.79 122.2 133.4 148.7 158.1 136.6 46.1 59.8 164.4 41.4 116.8 41.1 162.2 164.2 157.8 ... 150.0 135.16598.61 Ni i 4.24 -0.93 27.1 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 26.36635.14 Ni i 4.42 -0.75 25.1 ... 23.9 ... ... 36.9 ... ... 45.6 23.9 21.6 ... ... ... 68.7 ... 18.46300.31 O i 0.00 -9.75 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...6363.79 O i 0.02 -10.25 45.8 54.1 40.5 42.5 27.0 33.2 23.4 41.1 48.6 46.6 ... 31.9 37.8 46.1 38.9 41.9 55.75526.82 Sc ii 1.77 0.03 119.7 110.3 110.1 91.8 111.6 119.5 96.0 107.6 107.9 115.4 115.7 96.4 122.6 109.3 112.3 131.2 105.66245.62 Sc ii 1.51 -0.97 79.9 67.0 77.7 71.9 65.9 66.6 68.4 61.3 69.2 ... 52.8 78.5 54.1 61.2 82.8 54.5 60.66309.90 Sc ii 1.50 -1.52 58.0 54.7 ... 34.7 37.1 ... 18.7 39.9 ... ... 30.6 ... 54.7 43.5 ... ... ...6604.60 Sc ii 1.36 -1.31 101.0 88.9 94.8 82.7 84.2 98.3 58.4 83.7 113.2 93.9 54.5 102.4 87.0 82.5 126.4 71.7 82.66125.03 Si i 5.62 -1.57 ... ... ... 31.7 ... ... ... 32.3 38.8 ... ... 56.2 ... ... ... 17.4 ...6142.48 Si i 5.62 -1.51 ... ... ... ... ... ... ... 26.0 ... ... 16.5 ... ... ... ... 16.9 ...6145.02 Si i 5.61 -1.37 15.4 22.7 ... ... ... ... ... 22.1 ... 21.2 ... ... ... ... 19.3 ... 21.76155.14 Si i 5.62 -0.80 47.6 55.8 39.1 42.8 27.6 45.5 ... 41.0 38.9 51.2 34.6 49.0 40.0 39.0 43.5 47.9 45.06237.33 Si i 5.61 -1.02 26.6 23.7 32.1 16.0 ... 30.2 ... 33.6 23.4 43.1 22.3 31.0 26.8 ... 21.6 31.4 ...6243.82 Si i 5.61 -1.27 ... 53.1 44.0 ... ... ... 22.5 28.0 33.8 46.9 ... ... ... 29.3 17.0 23.1 ...5490.16 Ti i 1.46 -0.93 132.3 158.9 148.0 98.6 116.9 97.2 50.8 87.2 113.2 143.2 32.3 152.3 123.0 132.5 106.5 92.5 117.35503.90 Ti i 2.58 -0.19 76.8 94.2 120.5 66.3 65.9 56.5 23.4 53.6 52.8 107.6 ... 103.1 86.6 67.2 53.2 40.6 71.76126.22 Ti i 1.07 -1.42 167.5 167.6 172.5 130.1 130.7 127.2 56.8 138.8 148.4 182.3 32.4 172.4 162.0 159.6 138.8 119.0 160.56220.50 Ti i 2.68 -0.14 ... ... ... ... ... ... 24.1 ... 56.6 94.0 ... 81.3 80.0 ... ... ... ...6258.10 Ti i 1.44 -0.35 223.9 231.1 232.3 195.2 160.5 181.7 83.5 165.8 209.4 258.2 72.4 243.1 216.5 223.1 181.3 138.1 204.86303.77 Ti i 1.44 -1.57 121.7 117.2 115.7 91.1 79.0 87.1 41.0 77.9 87.3 121.2 ... 115.6 96.2 100.3 87.9 68.3 97.96312.24 Ti i 1.46 -1.55 ... 114.7 107.9 ... ... 79.2 20.6 78.5 91.3 113.1 ... 100.9 ... ... 88.7 45.6 86.36336.10 Ti i 1.44 -1.74 103.0 101.2 115.0 90.8 66.1 63.6 26.0 77.7 90.6 120.9 25.3 97.7 90.7 89.7 73.6 68.5 103.86508.12 Ti i 1.43 -2.05 69.8 90.0 72.0 56.9 45.3 53.3 ... 38.1 71.5 94.2 ... 90.2 58.0 54.3 63.9 31.0 63.56556.08 Ti i 1.46 -1.07 ... ... ... ... 138.0 ... ... ... ... ... ... ... ... 139.8 ... ... ...6599.13 Ti i 0.90 -2.09 153.9 186.8 181.3 147.5 111.3 136.9 46.9 109.4 167.9 173.1 24.5 194.5 159.1 174.3 ... 78.9 164.26666.53 Ti i 1.46 -1.62 41.8 50.1 59.5 ... ... ... ... ... 39.7 39.2 ... ... 29.8 33.9 46.2 ... ...5418.77 Ti ii 1.58 -2.11 111.2 97.0 96.3 99.6 115.3 93.2 97.6 116.2 83.3 109.4 122.1 92.7 109.8 102.2 83.0 110.0 96.56219.94 Ti ii 2.06 -2.82 ... ... 28.0 31.3 28.1 ... 21.6 48.8 ... 33.9 27.4 ... 40.7 ... ... 21.8 ...6559.58 Ti ii 2.05 -2.02 67.8 ... ... ... 72.9 ... ... ... ... ... ... ... 57.7 61.8 ... ... ...6606.95 Ti ii 2.06 -2.79 43.8 40.7 46.2 24.6 34.4 47.6 39.3 34.1 59.6 49.2 31.7 49.9 33.7 ... 88.0 27.1 31.16680.13 Ti ii 3.09 -1.86 23.1 29.5 33.0 17.5 27.9 ... 26.5 41.5 29.1 21.4 26.1 36.4 25.1 17.9 36.6 30.7 19.96119.53 V i 1.06 -0.32 ... ... ... ... ... ... ... ... ... 97.4 ... ... ... ... 74.5 ... 79.16128.33 V i 1.05 -2.30 ... ... ... ... 18.8 ... ... ... ... 30.5 ... ... ... ... ... ... ...6135.37 V i 1.05 -0.75 139.1 151.3 ... 88.9 103.1 102.4 23.7 97.8 113.4 145.4 18.9 145.0 114.8 116.7 109.2 53.5 116.86150.15 V i 0.30 -1.79 166.8 178.6 152.7 148.5 140.0 128.8 39.7 123.6 147.8 181.6 23.9 181.2 170.7 172.2 148.8 95.5 150.06199.19 V i 0.29 -1.29 222.5 220.2 219.8 188.6 170.3 182.9 80.9 185.5 188.0 224.1 44.0 216.4 206.0 200.3 177.1 122.5 176.86216.36 V i 0.28 -0.81 198.5 213.6 215.0 183.8 172.5 178.3 76.9 161.4 197.1 220.1 46.5 199.8 186.6 200.0 187.5 164.8 221.46224.51 V i 0.29 -2.01 146.9 157.1 166.9 131.9 119.3 125.7 31.5 117.9 141.9 170.8 ... 156.3 142.7 150.1 119.8 85.1 153.56233.20 V i 0.28 -2.07 106.0 107.5 117.6 111.4 80.0 87.4 ... 73.4 107.9 130.9 ... 130.1 109.1 119.3 95.1 54.2 113.56243.11 V i 0.30 -0.98 ... ... ... ... 261.0 ... 110.6 244.7 ... ... 47.0 ... ... ... ... 177.2 ...6251.82 V i 0.29 -1.30 171.7 166.0 160.4 128.5 137.9 143.1 73.3 139.6 163.9 182.3 29.4 172.7 168.1 155.8 141.1 94.5 174.56274.66 V i 0.27 -1.67 ... ... 138.9 ... ... 101.9 47.2 ... 135.5 151.8 ... 164.7 ... ... 120.0 77.7 128.96357.29 V i 1.85 -0.91 27.9 60.1 30.6 21.5 ... ... ... 30.3 27.8 55.9 ... 50.7 35.7 18.3 40.1 24.1 38.06452.32 V i 1.19 -1.21 97.3 128.1 140.8 99.8 57.9 38.1 15.2 100.2 51.9 119.0 ... 144.9 93.7 96.7 ... 48.1 97.56504.19 V i 1.18 -1.23 77.0 94.2 69.1 46.0 52.0 41.9 19.2 ... 62.7 94.4 ... 75.4 68.1 53.1 86.1 27.5 54.16531.41 V i 1.22 -0.84 84.0 ... ... ... 56.8 ... ... ... ... ... ... ... 62.7 79.8 ... ... ...5402.78 Y ii 1.84 -0.51 43.1 62.4 64.5 33.3 44.3 ... 27.9 50.5 32.5 63.1 ... 61.5 41.8 27.8 64.1 ... 51.86362.35 Zn i 5.80 0.14 ... 30.6 ... ... ... ... 21.2 ... ... ... ... ... ... ... 21.5 ... 19.86127.48 Zr i 0.15 -1.06 103.2 112.5 110.0 82.8 58.2 71.0 15.3 72.1 99.1 125.1 ... 121.8 106.2 95.2 87.5 53.4 108.86140.46 Zr i 0.52 -1.41 26.0 32.5 39.3 ... 19.7 22.9 ... ... 19.6 47.1 ... 45.9 27.2 28.0 ... 23.2 ...6143.18 Zr i 0.07 -1.10 110.6 124.4 114.9 82.2 77.8 55.6 18.4 62.9 95.0 150.3 18.9 110.1 114.0 101.3 89.8 45.5 115.56192.95 Zr i 0.54 -2.07 15.5 35.9 21.1 ... ... ... ... ... 24.0 29.0 ... 18.2 ... 16.9 ... ... 21.3

136chapter

6:H

Rspectroscopic

studyof

FornaxF

ieldStars

Table 6.A3: Abundance ratio of the Fornax Field stars, where we list the abundance, its associated error (σ) and the number of lines used (n). The quoted erroris the error on [element/H], not [X/Fe]. Part 1

Star [Fe i/H] σ n [Fe ii/H] σ (n) [Ba ii/Fe] σ n [Ca i/Fe] σ (n) [Cr i/Fe] σ n [Eu ii/Fe] σ (n) [La ii/Fe] σ n [Mg i/Fe] σ (n)BL038 -0.88 ± 0.07 (43) -0.70 ± 0.09 (4) 0.80 ± 0.16 (2) -0.25 ± 0.11 (8) -0.31 ± 0.23 (1) 0.62 ± 0.23 (1) 0.43 ± 0.16 (2) 0.03 ± 0.16 (2)BL045 -1.09 ± 0.05 (48) -0.78 ± 0.08 (5) -0.05 ± 0.21 (2) -0.33 ± 0.06 (8) -0.36 ± 0.16 (1) 0.19 ± 0.16 (1) 0.17 ± 0.16 (1) 0.02 ± 0.18 (2)BL052 -1.02 ± 0.06 (37) -0.52 ± 0.25 (2) 0.57 ± 0.39 (2) -0.42 ± 0.12 (8) -0.51 ± 0.34 (1) 0.67 ± 0.34 (1) 0.25 ± 0.34 (1) 0.07 ± 0.24 (2)BL065 -1.43 ± 0.07 (42) -1.06 ± 0.17 (5) 0.19 ± 0.13 (2) -0.09 ± 0.12 (9) -0.51 ± 0.18 (1) 0.35 ± 0.18 (1) ... ± ... (0) 0.00 ± 0.18 (1)BL076 -0.85 ± 0.06 (43) -0.43 ± 0.14 (4) 0.49 ± 0.16 (2) -0.21 ± 0.06 (9) -0.42 ± 0.19 (1) 0.34 ± 0.19 (1) 0.21 ± 0.19 (1) -0.15 ± 0.13 (2)BL077 -0.79 ± 0.06 (40) -0.38 ± 0.24 (4) 0.29 ± 0.13 (2) -0.46 ± 0.06 (8) -0.62 ± 0.15 (1) 0.29 ± 0.15 (1) -0.01 ± 0.11 (2) -0.27 ± 0.09 (3)BL079 -0.52 ± 0.08 (39) -0.16 ± 0.42 (2) 0.98 ± 0.44 (1) -0.26 ± 0.17 (7) -0.30 ± 0.44 (1) 0.45 ± 0.44 (1) 0.44 ± 0.31 (2) -0.06 ± 0.31 (2)BL081 -0.62 ± 0.06 (40) -0.44 ± 0.20 (4) 0.91 ± 0.13 (2) -0.34 ± 0.07 (9) -0.52 ± 0.19 (1) 0.26 ± 0.19 (1) 0.36 ± 0.13 (2) -0.32 ± 0.13 (2)BL084 -0.82 ± 0.05 (39) -0.59 ± 0.10 (4) 0.66 ± 0.21 (2) -0.30 ± 0.06 (9) -0.57 ± 0.19 (1) 0.44 ± 0.19 (1) 0.24 ± 0.13 (2) -0.12 ± 0.19 (1)BL085 -2.58 ± 0.05 (31) -1.99 ± 0.10 (4) 0.12 ± 0.11 (2) 0.33 ± 0.06 (9) ... ± ... (0) ... ± ... (0) 0.75 ± 0.15 (1) 0.35 ± 0.15 (1)BL091 -0.96 ± 0.06 (46) -0.60 ± 0.21 (5) 0.56 ± 0.12 (2) -0.15 ± 0.07 (7) -0.34 ± 0.17 (1) 0.52 ± 0.17 (1) 0.32 ± 0.12 (2) -0.04 ± 0.12 (2)BL092 -0.95 ± 0.08 (45) -0.57 ± 0.10 (3) 0.66 ± 0.16 (2) -0.23 ± 0.08 (9) -0.45 ± 0.23 (1) 0.53 ± 0.23 (1) 0.30 ± 0.16 (2) 0.04 ± 0.13 (3)BL096 -0.75 ± 0.08 (42) -0.95 ± 0.18 (4) 0.70 ± 0.18 (2) -0.18 ± 0.10 (6) -0.53 ± 0.25 (1) 0.53 ± 0.25 (1) 0.36 ± 0.25 (1) 0.02 ± 0.14 (3)BL097 -0.92 ± 0.05 (43) -0.53 ± 0.24 (5) 0.56 ± 0.20 (2) -0.18 ± 0.09 (9) -0.54 ± 0.17 (1) 0.48 ± 0.17 (1) 0.25 ± 0.12 (2) 0.06 ± 0.17 (1)BL100 -0.93 ± 0.06 (48) -0.67 ± 0.13 (5) 0.20 ± 0.25 (2) -0.48 ± 0.08 (8) -0.50 ± 0.19 (1) ... ± ... (0) 0.15 ± 0.13 (2) 0.00 ± 0.17 (3)BL104 -0.96 ± 0.06 (43) -0.33 ± 0.23 (4) 0.58 ± 0.15 (2) -0.20 ± 0.07 (9) -0.75 ± 0.21 (1) 0.64 ± 0.21 (1) 0.45 ± 0.21 (1) -0.01 ± 0.15 (2)BL113 -0.75 ± 0.06 (42) -0.65 ± 0.10 (5) 0.94 ± 0.21 (1) -0.08 ± 0.09 (9) -0.42 ± 0.21 (1) 0.45 ± 0.21 (1) 0.40 ± 0.15 (2) -0.03 ± 0.21 (1)BL115 -1.44 ± 0.06 (40) -0.85 ± 0.12 (4) 0.07 ± 0.11 (2) -0.28 ± 0.06 (8) -0.55 ± 0.16 (1) 0.12 ± 0.16 (1) 0.11 ± 0.11 (2) 0.01 ± 0.13 (3)BL123 -0.97 ± 0.05 (41) -0.78 ± 0.07 (4) 0.40 ± 0.12 (2) -0.30 ± 0.06 (8) -0.40 ± 0.17 (1) 0.40 ± 0.17 (1) 0.38 ± 0.12 (2) -0.04 ± 0.12 (2)BL125 -0.73 ± 0.07 (44) -0.61 ± 0.13 (3) 0.87 ± 0.18 (2) -0.21 ± 0.09 (8) -0.57 ± 0.25 (1) 0.51 ± 0.25 (1) 0.58 ± 0.25 (2) 0.02 ± 0.18 (2)BL132 -0.85 ± 0.06 (43) -0.61 ± 0.12 (4) 0.23 ± 0.21 (2) -0.45 ± 0.07 (8) -0.50 ± 0.20 (1) 0.24 ± 0.20 (1) 0.01 ± 0.14 (2) -0.13 ± 0.12 (3)BL135 -0.95 ± 0.07 (39) -0.69 ± 0.17 (4) 0.62 ± 0.15 (2) -0.30 ± 0.08 (7) -0.41 ± 0.21 (1) 0.47 ± 0.21 (1) 0.29 ± 0.15 (2) -0.02 ± 0.21 (1)BL138 -1.01 ± 0.06 (43) -0.58 ± 0.13 (5) 0.47 ± 0.21 (2) -0.18 ± 0.09 (9) -0.47 ± 0.17 (1) 0.55 ± 0.17 (1) 0.40 ± 0.12 (2) -0.08 ± 0.10 (3)BL140 -0.86 ± 0.07 (41) 0.22 ± 0.27 (5) 0.69 ± 0.13 (2) -0.30 ± 0.06 (9) -0.39 ± 0.18 (1) 0.62 ± 0.18 (1) 0.45 ± 0.13 (2) 0.04 ± 0.15 (3)BL141 -0.82 ± 0.06 (46) -0.42 ± 0.12 (5) 0.28 ± 0.15 (2) -0.33 ± 0.08 (7) -0.47 ± 0.21 (1) 0.39 ± 0.21 (1) 0.28 ± 0.15 (2) -0.11 ± 0.15 (2)BL146 -0.92 ± 0.06 (39) -0.64 ± 0.07 (5) 0.58 ± 0.20 (2) -0.36 ± 0.10 (8) -0.40 ± 0.16 (1) 0.41 ± 0.16 (1) 0.49 ± 0.11 (2) -0.14 ± 0.09 (3)BL147 -1.37 ± 0.07 (47) -0.87 ± 0.10 (4) 1.27 ± 0.23 (2) -0.30 ± 0.11 (9) -0.80 ± 0.32 (1) 1.38 ± 0.32 (1) 1.09 ± 0.23 (2) 0.43 ± 0.24 (3)BL148 -0.63 ± 0.11 (37) -0.56 ± 0.15 (4) 0.09 ± 0.33 (1) -0.29 ± 0.14 (8) -0.23 ± 0.33 (1) 0.35 ± 0.33 (1) 0.54 ± 0.23 (2) -0.13 ± 0.19 (3)BL149 -0.91 ± 0.07 (45) -0.83 ± 0.17 (5) 0.57 ± 0.18 (2) -0.30 ± 0.09 (8) -0.52 ± 0.25 (1) 0.43 ± 0.25 (1) 0.31 ± 0.18 (2) -0.06 ± 0.18 (2)BL150 -0.83 ± 0.06 (42) -0.41 ± 0.15 (4) 0.36 ± 0.16 (2) -0.18 ± 0.08 (7) -0.38 ± 0.22 (1) 0.50 ± 0.22 (1) 0.37 ± 0.16 (2) 0.01 ± 0.16 (2)BL151 -0.86 ± 0.07 (45) -0.60 ± 0.10 (3) 0.74 ± 0.13 (2) -0.25 ± 0.07 (8) -0.33 ± 0.18 (1) 0.46 ± 0.18 (1) 0.42 ± 0.13 (2) -0.06 ± 0.10 (3)BL155 -0.75 ± 0.07 (40) -0.52 ± 0.21 (3) 0.76 ± 0.20 (2) -0.17 ± 0.08 (9) -0.40 ± 0.22 (1) 0.39 ± 0.22 (1) 0.34 ± 0.16 (2) -0.03 ± 0.31 (2)BL156 -1.13 ± 0.07 (45) -0.62 ± 0.19 (5) 0.51 ± 0.13 (2) -0.28 ± 0.07 (7) -0.44 ± 0.19 (1) 0.65 ± 0.19 (1) 0.54 ± 0.13 (2) -0.01 ± 0.13 (2)BL158 -0.87 ± 0.07 (43) -0.40 ± 0.15 (4) 0.95 ± 0.23 (1) -0.38 ± 0.08 (9) -0.49 ± 0.23 (1) 0.64 ± 0.23 (1) 0.58 ± 0.16 (2) 0.06 ± 0.13 (3)BL160 -0.87 ± 0.06 (37) -0.52 ± 0.08 (5) 0.55 ± 0.12 (2) -0.31 ± 0.08 (9) -0.60 ± 0.15 (1) 0.47 ± 0.15 (1) 0.46 ± 0.15 (1) -0.20 ± 0.11 (2)BL163 -0.74 ± 0.07 (41) -0.52 ± 0.14 (4) 0.95 ± 0.20 (1) -0.08 ± 0.07 (8) -0.23 ± 0.20 (1) 0.51 ± 0.20 (1) 0.62 ± 0.14 (2) -0.16 ± 0.14 (2)BL166 -0.89 ± 0.07 (42) -0.54 ± 0.24 (4) 0.74 ± 0.16 (2) -0.23 ± 0.08 (7) -0.59 ± 0.22 (1) 0.42 ± 0.22 (1) 0.36 ± 0.16 (2) -0.07 ± 0.13 (3)BL168 -0.88 ± 0.08 (41) -0.57 ± 0.20 (4) 0.55 ± 0.13 (2) -0.12 ± 0.07 (7) -0.30 ± 0.18 (1) 0.54 ± 0.18 (1) 0.53 ± 0.13 (2) -0.11 ± 0.13 (2)BL171 -0.90 ± 0.06 (40) -0.19 ± 0.14 (4) 0.18 ± 0.13 (2) -0.26 ± 0.07 (7) -0.28 ± 0.18 (1) 0.44 ± 0.18 (1) 0.33 ± 0.13 (2) -0.03 ± 0.11 (3)BL173 -0.78 ± 0.07 (41) -0.55 ± 0.14 (5) 0.54 ± 0.15 (2) -0.36 ± 0.08 (8) -0.54 ± 0.21 (1) 0.40 ± 0.21 (1) 0.47 ± 0.15 (2) -0.12 ± 0.12 (3)BL180 -0.91 ± 0.07 (38) -0.68 ± 0.06 (1) 1.22 ± 0.34 (1) -0.25 ± 0.12 (8) -0.46 ± 0.34 (1) 0.58 ± 0.34 (1) 0.76 ± 0.24 (2) -0.08 ± 0.24 (2)BL185 -0.71 ± 0.07 (39) -0.78 ± 0.14 (4) 0.96 ± 0.16 (2) -0.06 ± 0.10 (9) -0.15 ± 0.22 (1) 0.37 ± 0.22 (1) 0.43 ± 0.18 (2) -0.07 ± 0.13 (3)


6.A:

Largetables

137Star [Fe i/H] σ n [Fe ii/H] σ (n) [Ba ii/Fe] σ n [Ca i/Fe] σ (n) [Cr i/Fe] σ n [Eu ii/Fe] σ (n) [La ii/Fe] σ n [Mg i/Fe] σ (n)BL190 -0.79 ± 0.05 (43) -0.60 ± 0.12 (4) 0.06 ± 0.13 (2) -0.43 ± 0.06 (8) -0.67 ± 0.18 (1) 0.39 ± 0.18 (1) 0.01 ± 0.15 (2) -0.32 ± 0.13 (2)BL195 -0.97 ± 0.06 (48) -0.77 ± 0.09 (4) 0.17 ± 0.25 (2) -0.20 ± 0.08 (8) -0.43 ± 0.19 (1) 0.43 ± 0.19 (1) 0.10 ± 0.19 (1) -0.07 ± 0.13 (2)BL196 -1.02 ± 0.06 (43) -0.76 ± 0.10 (5) 0.22 ± 0.13 (2) -0.34 ± 0.06 (9) -0.40 ± 0.18 (1) 0.50 ± 0.18 (1) 0.40 ± 0.13 (2) -0.22 ± 0.18 (1)BL197 -0.89 ± 0.07 (39) -0.79 ± 0.13 (2) 0.79 ± 0.30 (2) -0.42 ± 0.10 (7) -0.57 ± 0.24 (1) 0.54 ± 0.24 (1) 0.37 ± 0.17 (2) 0.05 ± 0.16 (3)BL203 -0.83 ± 0.07 (43) -0.55 ± 0.12 (3) 0.58 ± 0.19 (2) -0.19 ± 0.09 (9) -0.47 ± 0.27 (1) 0.51 ± 0.27 (1) 0.36 ± 0.19 (2) -0.07 ± 0.16 (3)BL204 -1.00 ± 0.07 (33) -0.74 ± 0.16 (2) 0.46 ± 0.47 (2) -0.15 ± 0.15 (5) -0.30 ± 0.34 (1) 0.92 ± 0.34 (1) 0.53 ± 0.24 (2) -0.10 ± 0.34 (1)BL205 -0.69 ± 0.07 (43) -0.63 ± 0.08 (5) 0.76 ± 0.16 (2) -0.21 ± 0.12 (7) -0.42 ± 0.19 (1) 0.31 ± 0.19 (1) 0.41 ± 0.13 (2) -0.07 ± 0.13 (2)BL208 -0.66 ± 0.06 (43) -0.73 ± 0.06 (3) 0.86 ± 0.16 (2) -0.07 ± 0.10 (9) -0.56 ± 0.22 (1) 0.24 ± 0.22 (1) 0.44 ± 0.16 (2) -0.08 ± 0.13 (3)BL210 -0.76 ± 0.07 (42) -0.37 ± 0.22 (5) 0.90 ± 0.15 (2) -0.32 ± 0.07 (8) -0.53 ± 0.21 (1) 0.47 ± 0.21 (1) 0.52 ± 0.15 (2) 0.06 ± 0.18 (3)BL211 -0.67 ± 0.07 (43) -0.95 ± 0.14 (4) 0.98 ± 0.16 (2) -0.20 ± 0.07 (9) -0.21 ± 0.22 (1) 0.40 ± 0.22 (1) 0.45 ± 0.16 (2) -0.15 ± 0.13 (3)BL213 -0.87 ± 0.06 (42) -0.63 ± 0.10 (5) 0.65 ± 0.15 (2) -0.32 ± 0.07 (9) -0.51 ± 0.21 (1) 0.52 ± 0.21 (1) 0.47 ± 0.15 (2) -0.19 ± 0.16 (2)BL216 -0.72 ± 0.07 (35) -0.93 ± 0.19 (2) 0.96 ± 0.18 (2) -0.18 ± 0.10 (7) -0.28 ± 0.25 (1) 0.55 ± 0.25 (1) 0.62 ± 0.18 (2) -0.03 ± 0.14 (3)BL218 -0.60 ± 0.09 (41) -0.60 ± 0.23 (3) 1.02 ± 0.26 (1) -0.26 ± 0.10 (7) -0.38 ± 0.26 (1) 0.29 ± 0.26 (1) 0.48 ± 0.18 (2) -0.09 ± 0.15 (3)BL221 -0.86 ± 0.06 (43) -0.75 ± 0.20 (5) 0.80 ± 0.15 (2) -0.27 ± 0.10 (9) -0.39 ± 0.21 (1) 0.68 ± 0.21 (1) 0.61 ± 0.15 (2) -0.11 ± 0.15 (2)BL227 -0.91 ± 0.08 (43) -0.65 ± 0.14 (3) 0.94 ± 0.18 (2) -0.21 ± 0.08 (9) -0.21 ± 0.25 (1) 0.70 ± 0.25 (1) 0.47 ± 0.18 (2) -0.05 ± 0.18 (2)BL228 -0.88 ± 0.05 (42) -0.49 ± 0.17 (4) -0.05 ± 0.13 (2) -0.45 ± 0.08 (9) -0.37 ± 0.16 (1) 0.31 ± 0.16 (1) 0.08 ± 0.11 (2) -0.16 ± 0.11 (2)BL229 -0.71 ± 0.06 (39) -0.48 ± 0.14 (4) 0.79 ± 0.17 (2) -0.25 ± 0.11 (8) -0.58 ± 0.24 (1) 0.48 ± 0.24 (1) 0.46 ± 0.17 (2) -0.16 ± 0.17 (2)BL233 -0.68 ± 0.07 (44) -0.46 ± 0.12 (4) 0.66 ± 0.16 (2) -0.47 ± 0.08 (8) -0.64 ± 0.22 (1) 0.36 ± 0.22 (1) 0.33 ± 0.16 (2) -0.13 ± 0.13 (3)BL239 -0.88 ± 0.07 (43) -0.67 ± 0.12 (4) 0.67 ± 0.21 (2) -0.21 ± 0.08 (8) -0.47 ± 0.22 (1) 0.26 ± 0.22 (1) 0.14 ± 0.16 (2) -0.15 ± 0.16 (2)BL242 -1.04 ± 0.07 (43) -0.73 ± 0.23 (3) 0.64 ± 0.16 (2) -0.45 ± 0.09 (9) -0.37 ± 0.21 (1) 0.47 ± 0.21 (1) 0.27 ± 0.16 (2) 0.15 ± 0.17 (3)BL247 -0.82 ± 0.09 (37) -0.61 ± 0.12 (5) 0.49 ± 0.11 (2) -0.30 ± 0.11 (9) -0.38 ± 0.15 (1) 0.47 ± 0.15 (1) 0.50 ± 0.11 (2) 0.10 ± 0.23 (2)BL250 -0.68 ± 0.09 (38) -0.51 ± 0.34 (4) ... ± ... (0) -0.06 ± 0.11 (7) -0.17 ± 0.26 (1) 0.26 ± 0.26 (1) 0.43 ± 0.18 (2) -0.03 ± 0.15 (3)BL253 -0.66 ± 0.06 (39) -0.61 ± 0.12 (4) 0.88 ± 0.14 (2) -0.32 ± 0.11 (8) -0.50 ± 0.20 (1) 0.42 ± 0.20 (1) 0.53 ± 0.14 (2) -0.20 ± 0.12 (3)BL257 -0.58 ± 0.07 (41) -1.09 ± 0.15 (3) 0.91 ± 0.24 (1) -0.30 ± 0.08 (8) -0.47 ± 0.24 (1) 0.32 ± 0.24 (1) 0.62 ± 0.17 (2) -0.11 ± 0.14 (3)BL258 -0.56 ± 0.08 (39) -0.53 ± 0.08 (4) 1.01 ± 0.23 (1) -0.16 ± 0.13 (9) -0.35 ± 0.23 (1) 0.49 ± 0.23 (1) 0.56 ± 0.16 (2) -0.29 ± 0.16 (2)BL260 -0.86 ± 0.06 (39) -0.36 ± 0.11 (3) 0.07 ± 0.13 (2) -0.30 ± 0.08 (9) -0.49 ± 0.19 (1) 0.32 ± 0.19 (1) 0.26 ± 0.13 (2) -0.27 ± 0.13 (2)BL261 -0.79 ± 0.09 (43) -0.37 ± 0.24 (4) 0.40 ± 0.15 (2) -0.36 ± 0.08 (8) -0.55 ± 0.21 (1) 0.46 ± 0.21 (1) 0.30 ± 0.15 (2) -0.13 ± 0.15 (2)BL262 -0.78 ± 0.07 (44) -0.44 ± 0.15 (5) 0.59 ± 0.25 (2) -0.19 ± 0.10 (7) -0.38 ± 0.27 (1) 0.33 ± 0.27 (1) 0.25 ± 0.19 (2) -0.01 ± 0.16 (3)BL266 -1.44 ± 0.06 (43) -1.16 ± 0.21 (4) 0.18 ± 0.14 (2) -0.18 ± 0.09 (10) -0.36 ± 0.20 (1) 0.45 ± 0.20 (1) 0.18 ± 0.14 (2) 0.04 ± 0.23 (2)BL267 -0.72 ± 0.08 (46) -0.64 ± 0.11 (5) 0.65 ± 0.14 (2) -0.16 ± 0.10 (9) -0.32 ± 0.20 (1) 0.37 ± 0.20 (1) 0.47 ± 0.14 (2) 0.00 ± 0.12 (3)BL269 -0.81 ± 0.08 (44) -0.47 ± 0.40 (2) 0.89 ± 0.20 (2) -0.22 ± 0.09 (9) -0.54 ± 0.28 (1) 0.64 ± 0.28 (1) 0.40 ± 0.20 (2) -0.06 ± 0.16 (3)BL278 -0.73 ± 0.07 (38) -0.87 ± 0.16 (3) ... ± ... (0) 0.04 ± 0.11 (8) -0.11 ± 0.25 (1) 0.40 ± 0.25 (1) 0.60 ± 0.18 (2) 0.00 ± 0.18 (2)BL279 -1.51 ± 0.08 (44) -1.31 ± 0.20 (3) 0.46 ± 0.15 (2) -0.03 ± 0.11 (9) -0.51 ± 0.21 (1) 0.64 ± 0.21 (1) 0.23 ± 0.15 (2) 0.18 ± 0.21 (1)BL295 -0.69 ± 0.08 (38) -0.25 ± 0.27 (3) ... ± ... (0) -0.07 ± 0.13 (8) -0.45 ± 0.25 (1) 0.44 ± 0.25 (1) 0.67 ± 0.18 (2) -0.04 ± 0.14 (3)BL300 -0.92 ± 0.07 (40) -0.60 ± 0.09 (4) 1.08 ± 0.13 (2) -0.16 ± 0.09 (7) -0.28 ± 0.19 (1) 0.46 ± 0.19 (1) 0.74 ± 0.13 (2) 0.18 ± 0.11 (3)BL304 -0.89 ± 0.06 (42) -0.68 ± 0.12 (4) 0.56 ± 0.14 (2) -0.34 ± 0.08 (8) -0.57 ± 0.20 (1) 0.38 ± 0.20 (1) 0.36 ± 0.20 (1) -0.09 ± 0.16 (2)BL311 -0.78 ± 0.06 (34) -1.00 ± 0.40 (2) 0.63 ± 0.34 (1) -0.25 ± 0.13 (7) -0.52 ± 0.34 (1) 0.85 ± 0.34 (1) 0.36 ± 0.24 (2) -0.17 ± 0.20 (3)BL315 -0.82 ± 0.10 (45) -0.62 ± 0.18 (4) 0.66 ± 0.18 (2) -0.34 ± 0.09 (6) -0.51 ± 0.21 (1) 0.69 ± 0.21 (1) 0.49 ± 0.15 (2) -0.01 ± 0.21 (1)BL323 -0.88 ± 0.06 (42) -0.69 ± 0.19 (3) 0.53 ± 0.18 (2) -0.56 ± 0.14 (8) -0.48 ± 0.26 (1) 0.43 ± 0.26 (1) 0.18 ± 0.18 (2) -0.18 ± 0.15 (3)

138chapter

6:H

Rspectroscopic

studyof

FornaxF

ieldStars

Table 6.A3: Abundance ratio of the Fornax Field stars, where we list the abundance, its associated error (σ) and the number of lines used (n). The quoted erroris the error on [element/H], not [X/Fe]. Part 2

Star [Na i/Fe] σ n [Nd ii/Fe] σ (n) [Ni i/Fe] σ n [O i/Fe] σ (n) [Si i/Fe] σ n [Ti i/Fe] σ (n) [Ti ii/Fe] σ n [Y ii/Fe] σ (n)BL038 -0.46 ± 0.16 (2) 0.22 ± 0.21 (3) -0.13 ± 0.06 (15) 0.34 ± 0.23 (1) 0.35 ± 0.17 (4) -0.16 ± 0.14 (7) 0.18 ± 0.13 (3) 0.10 ± 0.23 (1)BL045 ... ± ... (0) -0.06 ± 0.16 (1) -0.22 ± 0.07 (11) 0.22 ± 0.16 (1) 0.25 ± 0.24 (2) -0.23 ± 0.05 (10) 0.03 ± 0.12 (4) ... ± ... (0)BL052 -0.85 ± 0.34 (1) 0.22 ± 0.24 (2) -0.28 ± 0.10 (11) 0.28 ± 0.34 (1) 0.16 ± 0.34 (1) -0.37 ± 0.13 (7) -0.12 ± 0.49 (2) -0.16 ± 0.34 (1)BL065 ... ± ... (0) 0.49 ± 0.18 (1) -0.15 ± 0.08 (8) 0.56 ± 0.18 (1) 0.32 ± 0.21 (2) -0.18 ± 0.08 (7) 0.28 ± 0.13 (4) ... ± ... (0)BL076 -0.67 ± 0.13 (2) 0.20 ± 0.15 (3) -0.26 ± 0.08 (13) 0.20 ± 0.19 (1) -0.16 ± 0.11 (3) -0.21 ± 0.06 (9) 0.09 ± 0.18 (3) ... ± ... (0)BL077 -0.78 ± 0.15 (1) 0.23 ± 0.14 (3) -0.23 ± 0.06 (14) 0.25 ± 0.15 (1) -0.09 ± 0.11 (2) -0.38 ± 0.06 (10) -0.01 ± 0.15 (2) ... ± ... (0)BL079 -0.70 ± 0.31 (2) 0.36 ± 0.25 (3) -0.24 ± 0.12 (14) 0.12 ± 0.44 (1) -0.18 ± 0.31 (2) -0.05 ± 0.16 (8) -0.11 ± 0.45 (2) 0.08 ± 0.44 (1)BL081 -0.70 ± 0.13 (2) 0.21 ± 0.17 (3) -0.24 ± 0.06 (14) 0.14 ± 0.19 (1) -0.45 ± 0.13 (2) -0.19 ± 0.07 (10) -0.05 ± 0.16 (3) -0.34 ± 0.19 (1)BL084 ... ± ... (0) 0.21 ± 0.18 (3) -0.22 ± 0.07 (13) 0.05 ± 0.19 (1) -0.10 ± 0.24 (2) -0.38 ± 0.06 (9) 0.27 ± 0.23 (2) -0.40 ± 0.19 (1)BL085 ... ± ... (0) 1.53 ± 0.15 (1) 0.29 ± 0.15 (4) ... ± ... (0) 0.84 ± 0.15 (1) 0.56 ± 0.15 (1) 0.45 ± 0.26 (4) ... ± ... (0)BL091 -0.47 ± 0.12 (2) 0.42 ± 0.10 (3) -0.09 ± 0.07 (13) 0.10 ± 0.17 (1) -0.07 ± 0.18 (3) -0.25 ± 0.06 (9) 0.27 ± 0.21 (3) -0.31 ± 0.17 (1)BL092 -0.95 ± 0.16 (2) 0.30 ± 0.13 (3) -0.20 ± 0.09 (11) 0.14 ± 0.23 (1) -0.08 ± 0.12 (4) -0.23 ± 0.09 (9) 0.33 ± 0.20 (3) 0.20 ± 0.23 (1)BL096 -0.77 ± 0.18 (2) 0.27 ± 0.21 (3) -0.29 ± 0.08 (12) 0.34 ± 0.25 (1) -0.07 ± 0.18 (2) -0.19 ± 0.09 (8) 0.10 ± 0.14 (3) -0.32 ± 0.25 (1)BL097 -0.63 ± 0.12 (2) 0.30 ± 0.10 (3) -0.10 ± 0.07 (16) 0.28 ± 0.17 (1) -0.08 ± 0.13 (3) -0.19 ± 0.05 (10) 0.21 ± 0.11 (3) -0.06 ± 0.17 (1)BL100 ... ± ... (0) 0.30 ± 0.18 (2) -0.27 ± 0.07 (15) 0.16 ± 0.19 (1) -0.08 ± 0.16 (2) -0.41 ± 0.07 (8) 0.30 ± 0.21 (2) -0.41 ± 0.19 (1)BL104 -0.79 ± 0.21 (1) 0.34 ± 0.13 (3) -0.15 ± 0.07 (12) 0.37 ± 0.21 (1) 0.48 ± 0.23 (3) -0.13 ± 0.07 (9) 0.21 ± 0.12 (3) 0.05 ± 0.21 (1)BL113 -0.36 ± 0.15 (2) 0.25 ± 0.18 (3) -0.24 ± 0.06 (12) -0.05 ± 0.21 (1) 0.00 ± 0.10 (4) -0.18 ± 0.08 (7) 0.24 ± 0.18 (3) -0.19 ± 0.21 (1)BL115 ... ± ... (0) 0.48 ± 0.16 (1) -0.12 ± 0.07 (10) ... ± ... (0) 0.00 ± 0.16 (1) -0.25 ± 0.08 (8) 0.11 ± 0.14 (5) ... ± ... (0)BL123 -0.74 ± 0.17 (1) 0.34 ± 0.12 (3) -0.16 ± 0.05 (16) 0.30 ± 0.17 (1) 0.03 ± 0.10 (3) -0.19 ± 0.06 (9) -0.10 ± 0.21 (3) -0.09 ± 0.17 (1)BL125 -0.56 ± 0.21 (2) 0.39 ± 0.16 (3) -0.33 ± 0.06 (15) -0.05 ± 0.25 (1) -0.15 ± 0.12 (4) -0.09 ± 0.08 (9) 0.36 ± 0.20 (3) -0.01 ± 0.25 (1)BL132 -0.98 ± 0.14 (2) 0.14 ± 0.17 (3) -0.19 ± 0.07 (12) -0.11 ± 0.14 (2) 0.00 ± 0.15 (3) -0.38 ± 0.08 (7) 0.07 ± 0.14 (2) ... ± ... (0)BL135 -0.51 ± 0.15 (2) 0.37 ± 0.17 (3) -0.05 ± 0.08 (14) 0.43 ± 0.21 (1) 0.04 ± 0.18 (3) -0.02 ± 0.07 (9) 0.14 ± 0.15 (3) -0.10 ± 0.21 (1)BL138 ... ± ... (0) 0.23 ± 0.17 (3) -0.10 ± 0.08 (14) 0.44 ± 0.17 (1) 0.29 ± 0.17 (1) -0.10 ± 0.06 (8) 0.15 ± 0.14 (3) -0.15 ± 0.17 (1)BL140 -0.82 ± 0.18 (1) 0.32 ± 0.28 (2) -0.15 ± 0.08 (14) 0.13 ± 0.18 (1) -0.15 ± 0.10 (3) -0.24 ± 0.06 (9) 0.17 ± 0.13 (2) 0.06 ± 0.18 (1)BL141 -0.91 ± 0.18 (2) 0.25 ± 0.14 (3) -0.24 ± 0.06 (15) 0.03 ± 0.21 (1) -0.21 ± 0.16 (2) -0.38 ± 0.07 (9) -0.08 ± 0.09 (5) -0.45 ± 0.21 (1)BL146 -0.72 ± 0.11 (2) 0.27 ± 0.12 (3) -0.19 ± 0.06 (13) 0.28 ± 0.16 (1) 0.04 ± 0.08 (4) -0.16 ± 0.05 (9) -0.15 ± 0.16 (1) -0.08 ± 0.16 (1)BL147 -0.40 ± 0.32 (1) 0.88 ± 0.23 (3) -0.34 ± 0.11 (9) ... ± ... (0) 0.10 ± 0.18 (3) -0.32 ± 0.12 (7) 0.60 ± 0.27 (3) -0.04 ± 0.32 (1)BL148 -0.47 ± 0.23 (2) 0.73 ± 0.27 (3) -0.16 ± 0.09 (14) 0.44 ± 0.33 (1) 0.03 ± 0.23 (3) 0.06 ± 0.13 (6) -0.19 ± 0.20 (4) -0.19 ± 0.33 (1)BL149 -0.55 ± 0.25 (1) 0.05 ± 0.25 (1) -0.26 ± 0.08 (13) 0.16 ± 0.25 (1) 0.23 ± 0.17 (4) -0.40 ± 0.09 (8) 0.41 ± 0.14 (3) 0.03 ± 0.25 (1)BL150 -0.60 ± 0.16 (2) 0.41 ± 0.20 (3) -0.22 ± 0.11 (12) 0.05 ± 0.22 (1) 0.04 ± 0.13 (3) -0.15 ± 0.08 (8) 0.02 ± 0.17 (4) ... ± ... (0)BL151 -0.64 ± 0.13 (2) 0.31 ± 0.13 (3) -0.20 ± 0.06 (11) -0.04 ± 0.21 (2) -0.02 ± 0.20 (2) -0.12 ± 0.06 (9) 0.26 ± 0.12 (4) ... ± ... (0)BL155 -0.57 ± 0.16 (2) 0.64 ± 0.31 (3) -0.12 ± 0.09 (15) 0.27 ± 0.22 (1) -0.06 ± 0.13 (3) -0.07 ± 0.10 (10) 0.15 ± 0.13 (4) 0.16 ± 0.22 (1)BL156 ... ± ... (0) 0.35 ± 0.33 (2) -0.11 ± 0.08 (11) ... ± ... (0) 0.04 ± 0.12 (3) -0.43 ± 0.07 (8) 0.33 ± 0.20 (2) ... ± ... (0)BL158 -0.60 ± 0.16 (2) 0.51 ± 0.27 (3) -0.15 ± 0.08 (12) 0.17 ± 0.16 (2) 0.07 ± 0.12 (4) -0.40 ± 0.08 (11) 0.35 ± 0.12 (4) -0.13 ± 0.23 (1)BL160 -0.67 ± 0.15 (1) 0.46 ± 0.11 (3) -0.19 ± 0.06 (16) ... ± ... (0) 0.01 ± 0.09 (4) -0.11 ± 0.06 (8) 0.31 ± 0.16 (2) -0.05 ± 0.15 (1)BL163 -0.46 ± 0.14 (2) 0.50 ± 0.14 (3) -0.13 ± 0.06 (12) 0.28 ± 0.20 (1) -0.18 ± 0.14 (2) -0.05 ± 0.15 (8) 0.45 ± 0.21 (3) 0.01 ± 0.20 (1)BL166 -0.57 ± 0.16 (2) 0.49 ± 0.13 (3) -0.18 ± 0.07 (13) 0.17 ± 0.22 (1) -0.08 ± 0.18 (3) -0.11 ± 0.07 (10) 0.28 ± 0.23 (3) 0.00 ± 0.22 (1)BL168 -0.47 ± 0.18 (1) 0.29 ± 0.11 (3) -0.16 ± 0.06 (13) 0.17 ± 0.18 (1) -0.02 ± 0.20 (3) 0.00 ± 0.06 (8) 0.01 ± 0.10 (3) -0.32 ± 0.18 (1)BL171 ... ± ... (0) 0.30 ± 0.22 (3) -0.20 ± 0.07 (14) 0.30 ± 0.18 (1) ... ± ... (0) -0.23 ± 0.07 (8) 0.37 ± 0.17 (3) -0.25 ± 0.18 (1)BL173 -0.61 ± 0.15 (2) 0.46 ± 0.25 (3) -0.14 ± 0.08 (13) 0.65 ± 0.21 (1) 0.18 ± 0.21 (1) -0.19 ± 0.07 (9) 0.40 ± 0.19 (4) -0.21 ± 0.21 (1)BL180 -0.50 ± 0.24 (2) 0.65 ± 0.24 (2) -0.27 ± 0.09 (13) 0.17 ± 0.34 (1) 0.06 ± 0.20 (3) -0.18 ± 0.11 (9) -0.01 ± 0.22 (4) 0.22 ± 0.34 (1)BL185 -0.49 ± 0.16 (2) 0.41 ± 0.21 (2) -0.24 ± 0.09 (16) 0.13 ± 0.22 (1) -0.02 ± 0.11 (4) -0.08 ± 0.08 (10) 0.01 ± 0.11 (4) -0.12 ± 0.22 (1)


6.A:

Largetables

139Star [Na i/Fe] σ n [Nd ii/Fe] σ (n) [Ni i/Fe] σ n [O i/Fe] σ (n) [Si i/Fe] σ n [Ti i/Fe] σ (n) [Ti ii/Fe] σ n [Y ii/Fe] σ (n)BL190 ... ± ... (0) 0.27 ± 0.31 (2) -0.31 ± 0.07 (11) 0.09 ± 0.18 (1) -0.22 ± 0.16 (2) -0.40 ± 0.06 (9) 0.11 ± 0.14 (3) -0.50 ± 0.23 (1)BL195 -0.61 ± 0.13 (2) 0.03 ± 0.14 (3) -0.21 ± 0.07 (8) 0.10 ± 0.19 (1) -0.02 ± 0.16 (3) -0.17 ± 0.06 (9) 0.23 ± 0.15 (4) ... ± ... (0)BL196 -0.94 ± 0.18 (1) 0.31 ± 0.10 (3) -0.27 ± 0.06 (11) ... ± ... (0) 0.03 ± 0.18 (1) -0.28 ± 0.06 (9) -0.04 ± 0.14 (5) -0.05 ± 0.18 (1)BL197 -0.70 ± 0.17 (2) 0.24 ± 0.17 (2) -0.16 ± 0.07 (13) 0.19 ± 0.24 (1) -0.03 ± 0.17 (2) -0.30 ± 0.09 (7) 0.32 ± 0.28 (3) 0.13 ± 0.24 (1)BL203 -0.56 ± 0.19 (2) 0.30 ± 0.27 (2) -0.21 ± 0.08 (14) 0.13 ± 0.27 (1) -0.05 ± 0.24 (3) -0.19 ± 0.09 (9) 0.23 ± 0.32 (2) -0.28 ± 0.27 (1)BL204 -0.48 ± 0.34 (1) 0.60 ± 0.34 (1) -0.24 ± 0.11 (9) 0.40 ± 0.34 (1) 0.11 ± 0.34 (1) 0.01 ± 0.12 (8) 0.36 ± 0.44 (3) -0.20 ± 0.34 (1)BL205 -0.48 ± 0.13 (2) 0.42 ± 0.30 (2) -0.23 ± 0.06 (14) ... ± ... (0) -0.02 ± 0.13 (5) -0.20 ± 0.10 (9) 0.17 ± 0.12 (4) -0.14 ± 0.19 (1)BL208 -0.55 ± 0.16 (2) 0.36 ± 0.14 (3) -0.31 ± 0.06 (16) 0.15 ± 0.22 (1) -0.16 ± 0.14 (4) -0.13 ± 0.07 (9) 0.05 ± 0.29 (2) -0.54 ± 0.22 (1)BL210 -0.60 ± 0.15 (2) 0.50 ± 0.12 (3) -0.12 ± 0.07 (13) 0.16 ± 0.21 (1) -0.20 ± 0.19 (3) -0.19 ± 0.08 (10) 0.16 ± 0.13 (3) -0.08 ± 0.21 (1)BL211 -0.60 ± 0.16 (2) 0.32 ± 0.13 (3) -0.17 ± 0.07 (16) 0.15 ± 0.22 (1) 0.38 ± 0.26 (4) -0.03 ± 0.08 (8) 0.08 ± 0.19 (3) 0.04 ± 0.22 (1)BL213 -0.67 ± 0.21 (1) 0.36 ± 0.15 (2) -0.25 ± 0.06 (15) 0.29 ± 0.21 (1) 0.01 ± 0.10 (5) -0.14 ± 0.07 (9) -0.03 ± 0.17 (4) 0.15 ± 0.21 (1)BL216 -0.52 ± 0.18 (2) 0.36 ± 0.14 (3) -0.10 ± 0.08 (14) 0.24 ± 0.25 (1) -0.10 ± 0.14 (3) -0.02 ± 0.10 (6) -0.10 ± 0.20 (4) 0.06 ± 0.25 (1)BL218 -0.81 ± 0.18 (2) 0.26 ± 0.15 (3) -0.19 ± 0.08 (16) 0.09 ± 0.26 (1) -0.08 ± 0.26 (2) -0.15 ± 0.09 (8) -0.38 ± 0.13 (4) -0.05 ± 0.26 (1)BL221 -0.87 ± 0.21 (1) 0.51 ± 0.21 (2) -0.25 ± 0.07 (15) 0.16 ± 0.21 (1) -0.13 ± 0.15 (2) -0.26 ± 0.07 (9) 0.47 ± 0.18 (3) -0.06 ± 0.21 (1)BL227 -0.46 ± 0.18 (2) 0.44 ± 0.31 (3) -0.19 ± 0.07 (12) 0.29 ± 0.25 (1) 0.16 ± 0.10 (6) -0.12 ± 0.09 (8) 0.23 ± 0.26 (2) 0.19 ± 0.25 (1)BL228 -0.86 ± 0.11 (2) 0.41 ± 0.16 (1) -0.25 ± 0.07 (14) 0.22 ± 0.16 (1) -0.05 ± 0.12 (3) -0.19 ± 0.05 (9) 0.16 ± 0.20 (3) -0.30 ± 0.17 (1)BL229 -0.71 ± 0.17 (2) 0.40 ± 0.17 (3) -0.27 ± 0.09 (14) 0.24 ± 0.24 (1) -0.03 ± 0.10 (6) -0.24 ± 0.09 (7) 0.05 ± 0.12 (4) -0.46 ± 0.24 (1)BL233 -0.56 ± 0.16 (2) 0.25 ± 0.13 (3) -0.24 ± 0.08 (15) -0.01 ± 0.22 (1) -0.18 ± 0.12 (4) -0.21 ± 0.07 (9) 0.19 ± 0.13 (4) 0.11 ± 0.22 (1)BL239 -0.74 ± 0.22 (1) 0.14 ± 0.13 (3) -0.22 ± 0.08 (14) 0.11 ± 0.22 (1) -0.21 ± 0.13 (3) -0.29 ± 0.07 (9) -0.08 ± 0.27 (2) ... ± ... (0)BL242 -0.65 ± 0.21 (1) 0.32 ± 0.33 (2) -0.13 ± 0.06 (13) 0.19 ± 0.15 (2) 0.09 ± 0.15 (2) -0.23 ± 0.07 (8) 0.16 ± 0.24 (3) -0.07 ± 0.21 (1)BL247 -0.56 ± 0.11 (2) 0.94 ± 0.13 (2) -0.07 ± 0.08 (15) 0.23 ± 0.15 (1) 0.15 ± 0.18 (2) -0.23 ± 0.09 (10) 0.18 ± 0.17 (4) 0.43 ± 0.15 (1)BL250 -0.37 ± 0.18 (2) 0.39 ± 0.23 (3) -0.16 ± 0.10 (14) 0.05 ± 0.21 (2) 0.18 ± 0.18 (3) -0.04 ± 0.11 (9) -0.25 ± 0.18 (2) -0.12 ± 0.26 (1)BL253 -0.58 ± 0.14 (2) 0.36 ± 0.20 (3) -0.19 ± 0.07 (17) 0.15 ± 0.20 (1) -0.13 ± 0.11 (4) -0.21 ± 0.08 (7) -0.05 ± 0.12 (4) -0.17 ± 0.20 (1)BL257 -0.73 ± 0.17 (2) 0.31 ± 0.14 (3) -0.20 ± 0.07 (15) 0.17 ± 0.24 (1) -0.03 ± 0.14 (5) -0.19 ± 0.11 (9) 0.03 ± 0.23 (3) 0.04 ± 0.24 (1)BL258 -0.66 ± 0.16 (2) 0.50 ± 0.33 (3) -0.20 ± 0.07 (15) 0.03 ± 0.23 (1) -0.20 ± 0.18 (3) -0.22 ± 0.11 (9) 0.16 ± 0.18 (4) 0.10 ± 0.23 (1)BL260 -0.80 ± 0.19 (1) 0.36 ± 0.15 (3) -0.26 ± 0.07 (12) 0.25 ± 0.19 (1) 0.30 ± 0.28 (3) -0.19 ± 0.08 (8) -0.06 ± 0.10 (4) -0.21 ± 0.19 (1)BL261 -0.78 ± 0.15 (2) 0.45 ± 0.12 (3) -0.28 ± 0.10 (12) -0.02 ± 0.21 (1) -0.39 ± 0.21 (1) -0.27 ± 0.07 (9) 0.11 ± 0.09 (5) -0.03 ± 0.21 (1)BL262 -0.93 ± 0.27 (1) 0.26 ± 0.26 (2) -0.26 ± 0.08 (13) 0.11 ± 0.27 (1) -0.13 ± 0.16 (3) -0.12 ± 0.09 (9) 0.11 ± 0.30 (2) ... ± ... (0)BL266 ... ± ... (0) ... ± ... (0) -0.18 ± 0.07 (10) 0.38 ± 0.20 (1) 0.39 ± 0.20 (1) -0.13 ± 0.08 (9) 0.50 ± 0.12 (4) 0.07 ± 0.20 (1)BL267 -0.43 ± 0.14 (2) 0.12 ± 0.23 (3) -0.19 ± 0.09 (14) 0.18 ± 0.20 (1) -0.10 ± 0.10 (7) -0.16 ± 0.07 (8) 0.31 ± 0.13 (4) -0.06 ± 0.20 (1)BL269 -0.62 ± 0.20 (2) 0.27 ± 0.21 (3) -0.22 ± 0.10 (15) 0.27 ± 0.28 (1) 0.27 ± 0.20 (4) -0.13 ± 0.10 (10) 0.05 ± 0.31 (3) -0.25 ± 0.28 (1)BL278 -0.17 ± 0.18 (2) 0.50 ± 0.14 (3) -0.13 ± 0.07 (14) 0.16 ± 0.25 (1) -0.04 ± 0.12 (4) 0.16 ± 0.08 (9) 0.10 ± 0.12 (4) 0.10 ± 0.25 (1)BL279 0.04 ± 0.21 (1) 0.66 ± 0.21 (1) -0.03 ± 0.10 (12) ... ± ... (0) 0.34 ± 0.12 (3) -0.23 ± 0.13 (5) 0.67 ± 0.14 (4) ... ± ... (0)BL295 -0.50 ± 0.18 (2) 0.32 ± 0.21 (3) -0.19 ± 0.07 (15) -0.09 ± 0.25 (1) 0.04 ± 0.32 (3) -0.10 ± 0.08 (9) 0.21 ± 0.29 (3) 0.08 ± 0.25 (1)BL300 -0.47 ± 0.13 (2) 0.49 ± 0.11 (3) -0.22 ± 0.07 (14) 0.18 ± 0.19 (1) -0.02 ± 0.11 (3) -0.04 ± 0.07 (8) 0.09 ± 0.12 (5) -0.03 ± 0.19 (1)BL304 -0.79 ± 0.14 (2) 0.21 ± 0.12 (3) -0.19 ± 0.08 (14) 0.23 ± 0.20 (1) 0.11 ± 0.14 (2) -0.27 ± 0.10 (9) -0.12 ± 0.14 (3) -0.33 ± 0.20 (1)BL311 -0.68 ± 0.24 (2) 0.02 ± 0.34 (1) -0.16 ± 0.10 (12) 0.15 ± 0.34 (1) -0.12 ± 0.15 (5) -0.33 ± 0.11 (9) 0.47 ± 0.42 (3) 0.24 ± 0.34 (1)BL315 -0.85 ± 0.21 (1) 0.34 ± 0.21 (1) -0.32 ± 0.11 (15) 0.29 ± 0.21 (1) -0.04 ± 0.11 (6) -0.33 ± 0.07 (9) 0.33 ± 0.14 (4) ... ± ... (0)BL323 -0.76 ± 0.26 (1) 0.16 ± 0.26 (2) -0.23 ± 0.08 (15) 0.26 ± 0.26 (1) 0.12 ± 0.18 (2) -0.34 ± 0.09 (8) 0.07 ± 0.17 (3) 0.06 ± 0.26 (1)

Chapter 7Conclusions

The object of this thesis has been to determine detailed chemical abundances of indi-vidual stars in the nearby Fornax dwarf spheroidal galaxy, based on high resolution

observations with the VLT spectrographs UVES and FLAMES. Fornax, is one of thefew dwarf galaxies to have an extensive population of globular clusters (with 5), and ithas also had a complex field star formation history dominated by a star formation atintermediate ages. For this thesis, samples of individual stars were studied in both theglobular clusters and in the field star populations of Fornax.

7.1 New Data Reduction and Analysis Techniques

UVES is an Echelle spectrograph for which classical observations of single stars followedby careful data reduction and analysis star by star is possible. FLAMES is a power-ful new multiplexing spectrograph that was more challenging to use. It required newtechniques to be developed, from the preparation of the observing run to the last stepsof getting the stellar abundances. In the case of FLAMES the 100+ fibres that weretypically allocated to scientific targets made a high degree of automating very impor-tant. Another important aspect was the relatively low resolution (of the HR mode) ofGIRAFFE and also the limited wavelength coverage. This required careful adjustmentsand testing of the usual approach applied to UVES data.

An important achievement of this thesis is the method developed to analyse approxi-mately one hundred stellar spectra in a consistent and statistically robust manner, usingtools that are typically used on spectra with twice the resolution and larger wavelengthcoverage. This required bringing together several complex tasks, including accurate stel-lar atmospheric models, atomic data for the absorption lines, codes of line formation, EWmeasurements and signal extraction methods, all of which need to be properly includedand treated in order to obtain accurate results. We developed a pipeline that deliversstellar parameters and abundances in a controlled manner. This involved developingerror analysis and diagnostics to carefully test the robustness of the results.

142 chapter 7: Conclusions

7.2 The Fornax Globular ClustersThe Fornax dSph contains five globular clusters with a range of properties such as metal-licity, central concentration and Horizontal Branch structure. Using VLT/UVES we havetaken the first high resolution spectra of individual stars in the Fornax globular clusters.We obtained detailed chemical abundances for 9 individual RGB stars in 3 of the 5 For-nax globular clusters. This makes the Fornax globular clusters some of the very fewextra-galactic globular clusters that have been studied in this detail. From my resultsit is clear that Clusters 1, 2 and 3 were formed promptly and early in the history ofFornax dSph. They are over abundant in α-elements (O, Mg, Ca) at a similar level toGalactic clusters at identical [Fe/H], and the heavy element abundances (Y, Ba, Eu) inthe 3 clusters are compatible with dominant r-process enrichment. In addition, Cluster 1is found to be the most metal-poor globular cluster known, although the difference inmetallicity between Cluster 1 and M 92 or M 15 in the Milky Way is small.

Thus, despite their very different mass, morphology and global star formation history,the abundance patterns of individual stars in the Fornax GCs are almost identical tothose found in the Milky Way globular clusters. This extends to the ubiquitous deep-mixing patterns found among globular cluster members that were dected in two stars ofCluster 1 and Cluster 3, and the rare anomalies like the Eu-rich stars of Cluster 3 forwhich the only Galactic counterpart known to date is M15. This suggests that all globularclusters, regardless of their host galaxy, were formed with the same initial conditions attheir epoch of formation, namely the same pre- or self-enriched processes and identicalnucleosynthesis patterns.

7.3 Fornax Field starsThanks to the multi-fibre capability of FLAMES we have been able to make detailedabundance measurements of a large sample of 81 RGB stars in the central part of For-nax. This is a significant, even dramatic, improvement on the previous UVES sample of3 individual field stars. Our abundance ratios provide detailed information as to whatwere the chemical enrichment processes in Fornax, and how they differ from the MW.

We find that Fornax field stars exhibit unusually low α-element ratios, as well as Niand Na abundances. The [α/Fe] dependence on [Fe/H] is different from the Milky Way,meaning that there has been a different efficiency of chemical enrichment of the ISM.Fornax field stars are clearly predominantly enriched by s-process elements, showing thestrong role of (metal poor) AGB stars. This is clearly seen from the high [Ba/Y] ratios,compared to the Milky Way.

Our sample, which was randomly chosen from the entire breadth of the RGB, is domi-nated by a relatively young, relatively metal rich population (see Figure 7.1). This meansthat we have obtained the most detailed picture of the chemical enrichment of Fornaxduring the last ∼4 Gyrs. There is only one field star in our sample which appears tobe old and metal poor, and its properties are almost indistinguishable from the globularclusters in Fornax, and also from Galactic halo stars at the same [Fe/H]. Figure 7.1 showsthe main abundance results (alpha, Ba, and Eu) versus time. The ages were determined

7.3: Fornax Field stars 143

Figure 7.1: Here we show the result of our abundance analysis for alpha elements, andan s- and an r-process element compared to the ages of the individual stars observedin Fornax field star population. Representative error bars are shown on the bottom leftcorner of each panel.

from a colour-magnitude diagram finding the appropriate isochrone using the detailedspectroscopic abundances (Fe, and alpha). This allows us to determine how the differentabundances vary with time.

The [α/Fe] ratios were higher in the past and have gradually decreased towards morerecent times. This is a sign that SN Ia are becoming increasingly important with time,similar to what we were able to deduce from the [α/Fe] versus [Fe/H] plot of chapter 6.It is clear that the s-process abundances (e.g., [Ba/Fe]) show a slow increase in the con-tribution with increasing time (and [Fe/H]). This means that as the stellar population

144 chapter 7: Conclusions

becomes more metal rich there has been a steady rise in the Ba abundance, and the agesof the stars show that this rise began about 2-4 Gyr ago. The s-process is a much strongercontributor to the chemical evolution of Fornax than it is to the MW, or the SculptordSph. This suggests that stellar winds (e.g., from AGB stars) have played a uniquelyimportant role in the (recent, 2-4 Gyr ago) enrichment history of Fornax. There is nosuch trend in r-process, shown in the [Eu/Fe] panel of Figure 7.1. With the exception ofthe Eu-rich Cluster 3 points, all the observed stars in Fornax have a more or less constant[Eu/Fe].

Our detailed abundance studies confirm and deepen the difficulties found in earliermore limited surveys in understanding the role (if any) of dwarf galaxies, such as Fornaxin the build up of our Milky Way. These results also challenge our understanding of basicnucleosynthetic processes, with for example, ratios of [Ni/Fe] that are well below whatwas typically thought possible.

Further work on Fornax will be to investigate the different regions of this surprisinglycomplex dwarf galaxy in more detail, and specifically to increase our sample of highresolution abundances for metal poor stars in Fnx.

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Nederlandse samenvatting

Dit proefschrift draagt de titel “Chemische analyse van het Fornax dwerg sterrenstelsel”, enhet belangrĳkste doel van dit proefschrift is om te bepalen wat de chemische elementen

in de sterren van dit sterrenstelsel zĳn. Dit willen we achterhalen om zo te proberen deevolutie van dit sterrenstelsel te begrĳpen. Sterrenstelsels zĳn geen “statische” objecten; zebewegen, ze vormen sterren en ze kunnen interacties met andere sterrenstelsels hebben. Hetbestuderen van de sterren die een sterrenstelsel vormen kan ons informatie over het verledenvan het stelsel verschaffen. Sommige sterren kunnen zo oud zĳn als het sterrenstelsel zelf,terwĳl andere sterren veel jonger kunnen zĳn. Deze informatie kunnen we gebruiken om voorde gehele geschiedenis van de sterformatie in dit sterrenstelsel te bestuderen hoe de spectravan de sterren zĳn veranderd in de tĳd. De interpretatie van de resultaten is gebaseerd opde goed begrepen fysica van steratmosferen. Hier zal ik de achtergrond en de resultaten vandit proefschrift samenvatten.

Wat zĳn dwerg sterrenstelsels?Dwerg sterrenstelsels behoren in principe tot het meest simpele type stelsels dat we kennen.Door deze stelsels te bestuderen kunnen vele theorieën over de formatie en evolutie van sterrenen sterrenstelsels over een breed scala van omgevingen getest worden. Bolvormige dwergsterrenstelsels zĳn kleine, bĳna ronde sterrenstelsels die men voornamelĳk in de nabĳheidvan grotere stelsels, zoals de Melkweg, vindt. Normaal gesproken is er in deze stelsels geenstervorming en wordt er geen gas met deze stelsels geassociëerd. De abondantie ratio vande verschillende elementen in individuele sterren van verschillende leeftĳden verschaffen onseen gedetailleerde kĳk op de verschillende verrĳkkings-processen (Oftewel, supernovas, sterwinden) die op hun buurt ons begrip vergroten van de globale processen van de formatie enevolutie van het gehele sterrenstelsel.

De formatie van de elementenOp dit moment wordt aangenomen dat het Heelal is begonnen als een explosie, the Big Bang.In deze explosie werden waterstof, deuterium, helium en lithium gecreërd. Deze elementenworden dus gezien als de primaire elementen en alle andere elementen zĳn naderhand gevormddoor middel van nucleosynthesis in de sterren. Nucleosynthese is dus verantwoordelĳk voor

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bĳna alles wat we vandaag de dag op aarde zien. Dit proces werd in de jaren ’50 voor heteerst uitgelegd in werk van Fowler en Hoyle, wat uiteindelĳk leidde tot het B2FH (Burbidge,Burbidge, Fowler, & Hoyle 1957) artikel.

Het eerste, meest fundamentele proces om waterstof in zwaardere elementen te verande-ren, is waterstof verbranding. Dit is het proces waarbĳ waterstof in helium wordt veranderd.In lage massa sterren met een lage temperatuur in de kern gebeurt dit via de proton-protonketting en in de zwaardere sterren met hogere temperaturen via het invangen van protonendoor koolstof, stikstof en zuurstof atomen (in de CNO cyclus). Het volgende stadium ishelium verbranding. Daar waar helium zich verzameld in de kern van de ster, zal de kerngaan inkrimpen totdat de temperatuur en de dichtheid hoog genoeg zĳn voor een reactiewaarbĳ helium de brandstof is. Hierna volgt shell burning: koolstof verbranding, zuurstofverbranding, silicium verbranding. Deze processen kunnen uiteindelĳk elementen vormen zozwaar als 56Fe. Dit is het zwaarste element dat door kernfusie in de kern van een ster kanworden gevormd.

De belangrĳkste zwaardere elementen worden alpha elementen genoemd. Deze elementenhebben een kern die meervouden zĳn van helium. Dit zĳn de elementen O, Mg, Ca, Si enTi.Deze elementen worden voornamelĳk gevormd door alpha vangsten tĳdens de verschillen-de verbrandings fases van de massieve sterren. Deze elementen worden door supernova IIexplosies gedeponeerd in het interstellaire medium. En andere belangrĳke groep elementenis de ĳzerpiek. Deze groep wordt voornamelĳk geproduceerd door SN Ia explosies en ookdoor deze explosie terug gebracht in het ISM. Men verondersteld dat SN Ia explosie wordenveroorzaakt door de explosie van een witte dwerg die een dubbelster vormt met een minderzware progenitor ster. Deze explosies komen voornamelĳk voor ∼1 Gyr na de eerste ster-vormings episode, in tegenstelling tot SN II explosies die kort levende massieve sterren alsvoorlopers hebben (Dit kan minder dan ∼ 5− 10 Myrs zĳn).

Zwaardere elementen die voorbĳ de ĳzerpiek liggen worden gecreërd door het invangenvan neutronen. De twee belangrĳkste processen (vanuit een sterrenkundig oogpunt) hierbĳzĳn het s- en het r- proces. Als de neutronen flux niet hoog is dan krĳgen we het s-proces (ofslow-process). Hierbĳ zĳn de intervallen tussen het invangen van een neutron lang vergelekenmet de beta verval tĳdsschaal van een instabiele kern. De voorwaarden waaronder dit proceszich voordoet komen voornamelĳk voor in de omhulsels van thermisch pulserende AGB sterrenen zĳn het meest effectief in sterren met massa’s 3-5M�. Het r-proces (of rapid-process)manifesteert zich als de neutronen flux hoog genoeg is voor het snel invangen van neutronen.Deze omstandigheden verwacht men voornamelĳk in omgevingen zoals die gevormd woordendoor SNe II. Als de neutronen zo snel achter elkaar worden ingevangen kunnen ze accumulerenop instabiele kernen voordat deze kernen tĳd hebben voor een alpha of beta verval. De sterrendie verantwoordelĳk zĳn voor deze explosies zĳn massief. Deze sterren hebben daarom eenkorte levensduur en worden verondersteld de eerste sterren te zĳn die zware elementen in hetISM brengen.

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Abondanties in sterrenstelselsTĳdens het leven van een ster worden de abondanties van elementen gepreserveerd∗ op hetoppervlakte van de ster. Deze abondanties kunnen relatief makkelĳk worden achterhaald doorhet meten van absorptie lĳnen in spectra met een hoge resolutie. Hierdoor zĳn de abondantiesvan de elementen een belangrĳk instrument voor het begrĳpen van de oorsprong van een sterpopulatie geworden. De abondanties van verschillende elementen kunnen worden gemetenin sterren van verschillende leeftĳden. Dankzĳ de verschillende nucleosynthetische oorsprongvan de verschillende elementen kunnen we achterhalen wanneer welk verrĳkings proces domi-nant was in de geschiedenis van het sterrenstelsel. Niet geheel onverwacht, hebben de eerstestudies zich geconcentreerd op de Melkweg en pas onlangs zĳn vergelĳkbaar gedetailleerdestudies gedaan in andere sterrenstelsels zoals de Magellaanse wolken en meest recentelĳk dedichtbĳ gelegen bolvormige dwerg sterrenstelsels.

Deze studies suggereren dat de ster populaties van de satelliet stelsels die we vandaag dedag observeren geen significante bĳdrage kunnen leveren aan de ster populatie van de Melk-weg. Een mogelĳke uitzondering hierop is de buiten-halo van de de Melkweg. In vergelĳkingmet de Melkweg zĳn bolvormige dwerg stelsels simpelere stelsels. De meeste van deze stelselhebben een veel lagere hoeveelheid stervorming dan de Melkweg en al deze stelsels hebbeneen verschillend en uniek stervormings verleden. In dit proefschrift worden abondantie ratiobestudeerd in het Fornax dSph. Voor de eerste keer is het sample groot genoeg voor eengedetailleerde studie van de interne evolutie van deze stelsels en om het verschil met onzeMelkweg te kwantificeren.

Het kosmologisch belang van dwerg sterrenstelselsHet simpelste model voor de formatie van sterrenstelsels is dat alle sterrenstelsels wordengevormd door snel in elkaar te storten in het vroege Heelal. (het zogenoemde monolitisch inelkaar storten, of Eggen, Lynden-Bell, & Sandage 1962). Vervolgens evolueren deze stelselover de tĳd uitsluitend door hun gas in sterren te veranderen. Dit model verondersteld datde het overgrote deel van de massa in sterrenstelsels al bĳ het vormen van de stelsels op zĳnplaats was. Echter, dit model is herzien (e.g., Searle & Zinn 1978) tot een model dat veron-dersteld dat sterrenstelsels niet in een keer worden gevormd maar dat ze worden opgebouwduit kleinere segmenten. Tegelĳkertĳd met deze theorie kwam er het zeer succesvolle “colddark matter” (CDM) idee voor structuur formatie in het Heelal. CDM verondersteld dat dedonkere materie in een sterrenstelsel word opgebouwd door het continu invangen van kleineklonten om zo de stelsels en cluster die we vandaag de dag zien te vormen (e.g., White &Rees 1978; Navarro, Frenk, & White 1995).

Het aantal kleine satellieten rond een groter sterrenstelsel zoals het onze, lĳkt door CDMte worden overschat. Dit probleem staat bekent als het “missende dwerg probleem” (e.g.,Moore et al. 1999). De laatste jaren zĳn er vele zwakke satellieten rond om de Melkwegontdekt en deze ontdekkingen hebben onze ideeën over de lokale groep verandert. Dezeontdekkingen suggereren dat de bolvormige dwerg stelsels die we tot vandaag de dag hebbenbestudeerd nog maar het puntje van de ĳsberg zĳn. Deze stelsels zĳn de meest massieve∗ Behalve enkele lichte elementen die door intern mixen kunnen worden beïnvloed: Li, C, N.

154 Nederlandse samenvatting

satellieten van een grotere populatie van zwakkere satellieten met een lagere massa.

Wat hebben we geleerd?Een belangrĳk aspect van dit proefschrift is de pipeline die ontwikkeld is voor het analyserenvan een grote hoeveelheid stellaire spectra (∼ 100) op een consistente en statisch gezienrobuuste manier. Deze pipeline gebruikt gereedschappen die normaal gesproken worden ge-bruikt voor spectra met twee keer hogere resolutie en veel groter bereik in golflengte. Ditvereiste het samenbrengen van een stel complexe taken zoals, accurate ster modellen, atomai-re data voor de absorptie lĳnen, codes voor lĳn formatie, EW metingen en signaal extractiemethodes. Al deze taken moesten op de juiste wĳze in de pipeline worden geïncorpereerd enop de juiste wĳze worden gebruikt om nauwkeurige resultaten te krĳgen. De pipeline produ-ceert stellaire parameters en abondanties op een gecontroleerde manier. Dit vereiste ook deontwikkeling van een fouten analyse en diagnostiek zodat de robuustheid van de resultatenkon worden getest.

Het Fornax bolvormige dwerg sterrenstelsel heeft vĳf globular clusters (GCs) met eenbreedte aan eigenschappen. Door gebruik te maken van de VLT/UVES heb ik de eerstegedetailleerde chemische abondanties van negen individuele sterren in drie van deze GCs ver-kregen. Uit onze resultaten is het duidelĳk dat deze GCs direct en vroeg in de geschiedenisvan het Fornax dSph gevormd zĳn, netzoals de GCs in onze Melkweg. Ondanks het feit datdeze Fornax GCs in massa, morfologie en globale ster formatie geschiedenis zeer verschillen,zĳn de abondantie patronen van individuele sterren in deze GCs bĳna gelĳk aan de patronendie worden gevonden in GCs van de Melkweg. Ook abondantie patronen gevonden in sterrenzeer specifiek voor GCs (diep-mixen) en zeldzame anomalieën (europium-rĳk) zĳn bĳna iden-tiek. Dit suggereert dat sterren in GCs identiek zĳn ongeacht de grootte of het type stelselwaarmee zĳ geassocieerd worden.

Dankzĳ de multi-fibre mogelĳkheid van VLT/FLAMES heb ik gedetailleerde abondantiemetingen kunnen doen van 81 RGB sterren in het centrale deel van Fornax. Dit is een signifi-cante, zelfs dramatische, verbetering van het vorige UVES sample dat bestond uit slecht driesterren in het veld. Dit sample van Fornax veld sterren toont ongewoon lage [α/Fe] ratio,en de afhankelĳkheid met metallicity is verschillend van de Melkweg. Dit impliceerd eenverschillende efficiëntie in gas verrĳking. Fornax veld sterren zĳn overduidelĳk voornamelĳkverrĳkt door s-proces elementen met een hoge metallicity. Dit toont de belangrĳke rol van(metaal arme) AGB sterren aan. Ons sample word gedomineerd dooreen relatief jonge, me-taal rĳke ster populatie. Dit betekend dat we de chemische verrĳking van Fornax gedurendede laatste ∼4 Gyrs het meest gedetailleerd kunnen bekĳken. Ons sample bevat één veld steren deze lĳkt oud en metaal arm te zĳn. De abondantie eigenschappen van deze ster zĳn niette onderscheiden van de sterren in de globular clusters van Fornax.

Deze resultaten bevestigen en verdiepen de moeilĳkheden naar voren gebracht door eerderemeer gelimiteerde surveys in het begrĳpen van de rol die dit en vergelĳkbare stelsels hebbengespeeld bĳ het opbouwen van de Melkweg. Ook brengen deze resultaten een uitdaging onsbegrip van de nucleosynthese. De ratio voor [Ni/Fe] zĳn bĳvoorbeeld ver onder het niveaudat normaal gesproken mogelĳk werd geacht.

Résumé français

Cette thèse s’intitule « Analyse chimique de la galaxie naine Fornax » et son but principalest de tenter de comprendre l’évolution de cette dernière en déterminant les éléments

chimiques présents dans ses étoiles. Les galaxies ne sont pas des objets « statiques » ; ellesbougent, elles forment des étoiles et elles peuvent interagir avec d’autres galaxies. L’étude desétoiles composant une galaxie peut, en principe, nous informer sur son passé. Ainsi, certainesétoiles peuvent être aussi vieille que la galaxie elle-même ; d’autres peuvent être beaucoupplus jeunes. L’information sur le passé d’une galaxie peut servir à étudier la façon dont lesspectres stellaires ont varié pendant la période où les étoiles se sont formées dans dans laditegalaxie. L’interprétation des résultats repose sur la physique des atmosphères stellaires.

Que sont les galaxies naines ?En principe, les galaxies naines forment le type de galaxie le plus simple. On peut les étudierpour tester diverses théories sur la formation ainsi que l’évolution des étoiles et des galaxiesdans un éventail d’environnements. Les galaxies naines sphéroïdes sont petites, grosso modosphériques et sont habituellement trouvées dans les environs de galaxies plus larges, commela Voie lactée. Elles n’ont généralement pas d’étoiles en formation ni semblent-elles avoirde gaz qui leur est associé. Les ratios d’abondance de différents éléments dans des étoilesd’âges divers fournissent un aperçu détaillé des processus variés d’enrichissement chimique(supernovas, vents stellaires) qui à leur tour, améliorent notre compréhension des processusglobaux de la formation et de l’évolution des galaxies.

Formation des éléments chimiquesOn croit que l’Univers a débuté par une explosion, le Big Bang, où l’hydrogène, le deutérium,l’hélium et le lithium ont été créés. Ces derniers sont alors considérés comme des élémentsprimordiaux, et tous les autres éléments sont formés subséquemment dans les étoiles parnucléosynthèse. La nucléosynthèse stellaire est donc responsable de presque tout ce qui nousentoure aujourd’hui sur Terre. Ceci a d’abord été expliqué dans les années 1950 par le travailde Fowler et Hoyle, et finalement dans l’article des B2FH (Burbidge, Burbidge, Fowler, &Hoyle 1957).

156 Résumé français

Le tout premier procédé fondamental de conversion d’hydrogène en éléments plus lourdsest la combustion d’hydrogène. Il s’agit de convertir un noyau d’hydrogène en hélium, par lachaine proton-proton dans les étoiles de petites masses et dont le coeur a une basse tem-pérature et également par des captures de protons par des atomes de carbone, d’azote etd’oxygène (dans les cycles CNO) dans des étoiles plus massives et de températures plusélevées. L’étape suivante est la combustion de l’hélium. L’hélium s’accumule au coeur del’étoile et le coeur se contracte jusqu’à ce que la température ainsi que la densité augmenteassez pour produire une réaction dans laquelle l’hélium en est le carburant. Ensuite vient lacombustion des couches, où respectivement le carbone, l’oxygène et le silicium brulent. Ceprocédé peut produire des éléments aussi lourd que le 56Fe qui est l’élément le plus massifpouvant être formé par fusion au coeur d’une étoile.

Les éléments lourds les plus significatifs sont appelés les éléments alpha, avec des noyauxqui sont des multiples de l’He (O, Mg, Ca, Si, Ti). Les éléments alpha sont majoritairementsynthétisés par la capture d’alpha pendant les multiples phases de combustion dans les étoilesmassives et ils sont expulsés dans le milieu interstellaire (interstellar medium – ISM) par desexplosions de supernovas de type II (SN II). Un autre groupe important est formé par leséléments du pic du fer (iron peak), incluant le Fe lui-même. Ils sont produits essentiellementet expulsé dans le milieu interstellaire par des explosions de supernovas de type Ia (SN Ia),qu’on pense être dues à l’explosion d’une naine blanche dans un système binaire développéavec une étoile génitrice moins massive. Les SN Ia surviennent habituellement environ 1 mil-liard d’années après le premier épisode de formation d’étoile. Les SN II, quant à elles, ont desétoiles massives génitrices à vie courte, seulement 5-10 millions d’années.

Les éléments plus lourds, au-delà du pic du fer, sont créés par capture de neutrons oùles deux procédés les plus importants (dans un contexte astrophysique) sont les processus Set R. Le processus S (lent) survient lorsque le flux de neutrons est peu élevé, pour que lesintervalles entre les captures de neutrons sont longs par comparaison avec l’échelle de tempsde désintégration bêta d’un noyau instable. Ces conditions sont trouvées dans les enveloppesd’étoiles AGB (Asymptotic Giant Branch – branche asymptotiques des géantes) en pulsationthermique et elles sont plus efficaces dans des étoiles de 3-5 masses solaires. Le processus R(rapide) survient lorsqu’il y a un flux de neutrons suffisant pour permettre une capture rapidede neutrons. On croit que cela survient généralement dans des environnements similaires àceux produits par les SN II. Avec de telles captures rapides successives, les neutrons peuvents’accumuler sur un noyau instable avant même de subir une désintégration alpha ou bêta.Les étoiles responsables de ces explosions sont massives, donc elles ont une courte vie, et oncroit qu’elles sont les premiers objets fournissant des éléments lourds dans l’ISM.

Abondances dans les galaxiesParce que les abondances d’éléments sont préservées∗ à la surface d’une étoile pendant toutesa durée de vie, et parce qu’elles peuvent être (relativement) facilement mesurées à partir deraies d’absorbtion dans des spectres stellaires à haute résolution, les abonances sont devenuesun outil très important pour comprendre la genèse des populations stellaires. Les abondancesde divers éléments peuvent être mesurés dans des étoiles de différents âges et (en raison de∗ Sauf pour quelques éléments légers qui peuvent être affectés par mélange interne : Li, C, N.

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leur différente origine nucléosynthétique) elles nous permettent de déduire le procédé d’en-richissement qui a dominé à différentes époques de l’histoire de la galaxie. Les premièresétudes se sont évidemment concentrées sur la Voie lactée (VL), et ce n’est que relativementrécemment que des études détaillées similaires ont été faites dans d’autres galaxies, telles lesNuages de Magellan et, plus récemment, les galaxies naines sphéroïdales avoisinantes.

Ces études suggèrent que les populations stellaires des galaxies satellites que nous voyonsaujourd’hui ne peuvent avoir contribué de façon significative à la population stellaire de notreGalaxie, sauf peut-être pour le halo externe. Les galaxies naines sphéroïdales sont de simplessystèmes en comparaison avec la VL. La plupart d’entre elles ont généralement des taux deformation d’étoiles plus faibles et chacunes d’entre elles ont une histoire unique de formationstellaire. Dans cette thèse, un échantillon statistiquement significatif de ratios d’abondanceest étudié pour la première fois dans la galaxie naine sphéroïdale Fornax. Cette thèse faitune étude détaillée de l’évolution interne de ce système et compare les résultats trouvés avecnotre Galaxie.

Importance cosmologique des galaxies nainesLe modèle le plus simple de formation de galaxie stipule que toutes les galaxies se formentdans l’Univers primitif dans un scénario d’effondrement rapide (effondrement monolithique,Eggen, Lynden-Bell, & Sandage 1962). Au fil du temps, ces galaxies se développent ensuiteuniquement en changeant leur masse gazeuse en masse stellaire. Ce modèle indique que lamajorité de la matière de toutes les galaxies étaient en place à leur formation. Cependant, cemodèle simple a été surclassé (ex. Searle & Zinn 1978) par un modèle qui tient pour acquisque les galaxies ne se forment pas en un seul effondrement, mais plutôt qu’elles se constuisentavec le temps à partir de petits fragments. Cette théorie est venue en parallèle de la très po-pulaire vision de la formation de structure dans l’Univers, « matière sombre froide » (colddark matter – CDM), qui établit que le contenu en matière sombre d’une galaxie se construitpar l’accumulation continu de petits agglomérats, pour finalement former les galaxies et lesamas de galaxies que nous voyons aujourd’hui (ex. White & Rees 1978; Navarro, Frenk, &White 1995).

La CDM semblent surestimer le nombre de petites galaxies satellites autour de galaxiesplus larges (comme la nôtre) ; il s’agit d’une incohérence connue comme le « problème desnaines manquantes » (missing dwarf problem, ex. Moore et al. 1999). Cependant, de récentesdécouvertes au sujet de plusieurs satellites de faible luminosité présents autour de la VL de-puis quelques années changent notre point de vue du Groupe local, suggérant ainsi que lesgalaxies naines sphéroïdales étudiées jusqu’à maintenant ne sont que le bout de l’iceberg. Cesont les satellites les plus massifs d’une population plus large de satellites peu lumineux et deplus petites masses.

158 Résumé français

Qu’avons-nous appris ?Un aspet important de cette thèse est le système semi-automatique développé pour analy-ser un large nombre de spectres stellaires (∼ 100) de manière constante et statistiquementrobuste, en utilisant des outils qui sont généralement utilisés sur des spectres ayant unerésolution deux fois plus grande et une plus large couverture spectrale. Il a fallu joindre plu-sieurs tâches complexes, incluant des modèles d’atmosphères stellaires précis, des donnéesatomiques pour les raies d’absorbtion, des codes de formation de raies spectrales, des mé-thodes de mesure de largeur équivalente et d’extraction du signal. Toutes ces tâches doiventêtre incluses adéquatement et être traitées pour obtenir des résultats précis. Nous obtenonsdonc des paramètres stellaires et des abondances de façon contrôlée, en tenant compte ducalcul d’erreurs pour vérifier soigneusement la validité des résultats.

Fornax contient cinq amas globulaires (AG), ayant un éventail de propriétés. En utilisantle VLT/UVES, j’ai obtenu les premières abondances chimiques détaillées de neuf étoiles in-dividuelles dans trois de ses AG. À partir de nos résultats, il appert qu’ils ont été formésrapidement et tôt dans l’histoire de Fornax tout comme l’ont été ceux de la VL. Ainsi, malgréleur masse, leur morphologie et leur histoire globale de formation d’étoiles complètement dif-férentes, les modèles d’abondance des étoiles individuelles dans les AG de Fornax sont presqueidentiques que ceux trouvés dans les AG de la Voie lactée, incluant des modèles d’abondancespécifiques aux amas d’étoiles (brassage des couches profondes, deep-mixing) et de raresanomalies (surabondance en europium) également observés dans d’autres AG. Ceci suggèreque les étoiles dans les AG sont les mêmes sans égard pour le type ni pour la taille de lagalaxie d’où l’amas globulaire provient.

Grâce à la capacité multifibre de VLT/FLAMES, j’ai pu prendre des mesures détailléesd’abondance de 81 étoiles de la branche des géantes rouges (dans la partie centrale de For-nax). Ceci est une amélioration significative par rapport au précédent échantillon de troisétoiles individuelles du champ central provenant de UVES. Les étoiles de Fornax sont sous-abondantes en éléments alpha, et leur dépendance en métallicité est différente de la VL. Cecisuggère une efficacité différente d’enrichissement du gaz. À haute métallicité, les étoiles sontessentiellement enrichies en éléments du processus S, démontrant le rôle important des étoilesAGB pauvre en métaux. Notre échantillon est dominé par une population relativement jeuneet relativement riche en métaux. Nous avons donc obtenu un portrait détaillé de l’enrichisse-ment chimique de Fornax pour les dernières ∼4 milliards d’années. Il n’y a qu’une d’étoilesdans notre échantillon qui semble vieille et pauvre en métaux, et ses propriétés d’abondancesont presque indifférenciable de celles des AG de Fornax.

Ces résultats confirment les difficultés trouvées dans les précédentes études (plus limi-tées) quant à la compréhension du rôle (s’il y en a un) de la galaxie Fornax (ou d’une galaxiesimilaire) dans la formation de la VL. Ces résultats mettent également au défi notre compré-hension des processus nucléosynthétiques de base, avec, par exemple, des ratios de [Ni/Fe]qui sont nettement plus bas que ce qui est perçu généralement comme possibe.

Acknowledgements

There are many people and organisations I would like to thank for their help in the makingthis thesis. First, I would like to thank the Dutch society for giving me the opportunity tostudy in the Netherlands, and Groningen for being such a nice place to live. Second, theNetherlands Organisation for Scientific Research (NWO) for funding my research project, theKapteyn Astronomical Institute and the Leids Kerkhoven-Bosscha Fonds (LKBF) for fundingmy travel. And for the printing of this thesis, I’m grateful to the LKBF, the Central StudentAdministration Office and the Kapteyn Astronomical Institute.

Thank you Eline and Vanessa for supervising me during the last four years. I know thatthe last part was tough and that time was flying, but you kept pushing me in the right di-rection. Your help in reading, correcting and commenting my thesis was invaluable and I amgrateful for that. You have been patient with me, given me the chance to retry when I madea mistake, shown me the way when I was lost and always made sure that at least one of youcould help me when I was needing it. Thank you Mike, Kim, Matthew and Amina for yourcontributions to my work all over the years. And a special thank to Pascale: you were one ofthe most helpful person during my PhD for all the discussions we had together, always readyto give an advice, with your colourful style, direct, funny, and most of the time right. Youwere never afraid to say what had to be said and I’m glad I have been collaborating with you.

I’m also grateful to the reading committee of my thesis, Monique Spite, Piet van der Kruitand Jan Willem Pel, for their quick reading and insightful comments, helping to improve thisthesis in many aspects. Your comments and suggestions were greatly appreciated and willprobably help everyone who will my thesis. Many thanks to Bertrand Plez for making thestellar models we needed for Fornax; without your contribution, this work would not havebeen the same. Thanks to Marco Gullieuszik and Enrico Held for providing JHK photometryfor some of our targets, and to Serge Demers for his carbon star list.

Thank you Giuseppina for being my house mate for three years, we really had some funin this small cosy house on the Nieuwe Kĳk in’t Jatstraat! Thanks for helping me improvemy now famous pasta alla Bruno. All the dinners we organised, the movies night, your homemade liquors and the memorable parties and wine tasting nights we had (that would inevitablyend up in Storm for one last beer). I have some really funny memories about this house, likethat time when it was really late in the night and there was some noise outside and you...

160 Acknowledgements

Thank you Yang-Shyang for sharing your flat with me in the last six months of my stay inthe Netherlands. You were a great moral support in that difficult time, cooking for me muchmore that I cooked for you, supporting me everyday in many aspects and trying to make surethat my life was not too boring even when I was working 16 hours per day. You were veryunderstanding, always there for me and never asking anything in return for your generosity.

Thanks to César, Erik, Miguel and Wilfred for being my office mates. Thanks to Katarina,Miguel, Pablo and Fabrice for your moral support while going through the end of your thesisat the same time as me. Thank you Rien for all the political discussions we had (mostly)about Canada. Even though you were wrong most of the time, I enjoyed arguing with youon every occasions, and I will be missing this. Thanks to Bernard, Mark, Nigel, Saleem andScott for all the political, social and scientific discussions we had. Thanks to Paolo andSimona for painting my beloved moped in pink, that was too generous of you, I will try toreturn the favour one day. Thanks to Jackie for finding a house for me to live in on thefirst day I arrived in the Netherlands. Rense, Jelte, Bjorn, Teffie, Edo ×2, Peter, Erwin,Michiel, Chris, Christiaan, Dieter and Wilfred, you helped me on countless occasions whenI needed help to translate something in dutch, thank you. Teffie, thanks for making it allthe way to Paris to visit me, and for all the dinners and discussions we had. Peter, Jan,Giussepina, Mirjam, Yang-Shyang and Alvaro, thanks for your company in Storm. Dieter,thank you for your guitar lessons and advices, and give my thanks again to your father for hisnice historical tour of Gent. Thank you Patricia for your help with my postdoc applications.Thanks to Renzo for your advice on where to eat haring. Fabrice, Philippe and Ulrike, thanksfor inviting me for dinner on several occasions, it was nice to have someone to speak frenchin Groningen. Thanks to Scott and Danny for helping me with Python and Linux. Thankyou Eite for your quick response in getting a new desktop computer for me when my laptopfailed. Thank you Monty Python for everything you’ve done. Thanks to Jeffrey Lebowskifor being such a great character. Thank you Plume and Renaud for all the songs you wrote.Thanks to Pierre Falardeau for “Le temps des bouffons”. Thanks to Denis Drolet for beingdressed in brown. Thanks to J.R.R. Tolkien and Douglas Adams for writing excellent stories.Thanks to de Pintelier for the good selection of beers, and to Storm for the good musicand ambiance. Thanks to Images for the yearly French film festival. Thanks to Ole, Rense,Jelte, Teffie, Edo, Peter, Dieter, Wilfred, Paolo, Mirjam, Alvaro, Giussepina, Yang-Shyang,Martin, Matias, Facundo, Simona, Mercedes, Katarina, Elif, Seungyoup, Lodovico, Jeronimo,Emilio, César, Erik, Monica, Isabel, Léon, Patricia, Filippo, Philippe, Fabrice, Alicia, Elineand Andrew for all the dinners, parties, game nights and movie nights.

I also want to thank the Observatoire de Paris for welcoming me during my visits, and theCNRS for funding part of these visits. Special thanks to Vanessa, Jean-Noël, Selma, Pascaleand Aurélie for taking good care of me while I was in Paris. Thank you Martin, Yuyan,Martine, Jesse and Keylan for your hospitality in December and January, when I becamehomeless after leaving the Netherlands. Thanks to Peter Kamphuis and Théane Lavigne formaking the dutch and french summary for my thesis on such a short notice, 1 day! Thank youJesse for the nice cover page you made for my thesis. Thank you Sylviane for all your helpand financial advices. Thank you Céline, Raynald, Martine and Simon for always encouragingme, always being there for me and providing stability in my life for the past 30 years. – Merci!

Bruno Letarte – February 2007 – Pasadena

Documents

Chemical Analysis of the Fornax Dwarf Galaxy