Elemental abundances in the Milky Way stellar disk(s), bulge, and halo

New Astronomy Reviews 57 (2013) 80–99

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Elemental abundances in the Milky Way stellar disk(s), bulge, and halo Sofia Feltzing a,⇑, Masashi Chiba b a b

Lund Observatory, Department of Astronomy and Theoretical Physics, Box 43, SE-22100 Lund, Sweden Astronomical Institute, Tohoku University, Sendai 980-8578, Japan

a r t i c l e

i n f o

Article history: Available online 8 July 2013

a b s t r a c t We present a review of elemental abundances in the Milky Way stellar disk, bulge, and halo with a focus on data derived from high-resolution stellar spectra. These data are fundamental in disentangling the formation history and subsequent evolution of the Milky Way. Information from such data is still limited and confined to narrowly defined stellar samples. The astrometric Gaia satellite will soon be launched by the European Space Agency. Its final data set will revolutionize information on the motions of a billion stars in the Milky Way. This will be complemented by several ground-based observational campaigns, in particular spectroscopic follow-up to study elemental abundances in the stars in detail. Our review shows the very rich and intriguing picture built from rather small and local samples. The Gaia data deserve to be complemented by data of the same high quality that have been collected for the solar neighborhood. Ó 2013 Elsevier B.V. All rights reserved.

Contents 1.

2.

3. 4.

Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1.1. Chemical clocks . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1.2. Differential analysis. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . Characterizing the stellar disk(s). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.1. In the solar neighborhood. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.2. Moving groups in the solar neighborhood and the concept of chemical tagging . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.3. How well mixed is the solar neighborhood?. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.4. Super-metal-rich stars: visitors from the Galactic bulge? . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.5. Beyond the immediate surroundings . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.5.1. Old and young stellar components . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.5.2. Radial abundance gradients in other galaxies. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.5.3. Above the Galactic plane . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.6. Formation scenarios: a brief overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.6.1. Monolithic dissipative collapse . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.6.2. Multiple dissipative mergers at high redshift . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.6.3. Tidal debris of shredded satellites . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.6.4. Disk heating triggered by minor merger. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.6.5. Radial migration of disk stars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.6.6. Clumpy disk formation at high redshift . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.6.7. Galactic chemical evolution . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . Elemental abundance trends in the Galactic bulge . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . Characterizing the stellar halo. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.1. Multiple halos: a long history of tracing the evidence . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.2. Insight from high-resolution abundance studies. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.3. Testing the merger picture of halo formation: insight from abundance studies. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.4. Chemodynamical evolution of the stellar halo . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

⇑ Corresponding author. Tel.: +46 46 2227294. E-mail addresses: sofi[email protected] (S. Feltzing), [email protected] (M. Chiba). URLs: http://www.astro.lu.se/~sofia (S. Feltzing), http://www.astr.tohoku.ac.jp/~chiba (M. Chiba). 1387-6473/$ - see front matter Ó 2013 Elsevier B.V. All rights reserved. http://dx.doi.org/10.1016/j.newar.2013.06.001

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Outlook . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 95 Acknowledgments . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 95 References . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 95

1. Introduction

1.1. Chemical clocks

Galaxy formation and evolution represent a key topic in contemporary astrophysics. Understanding the formation and evolution of galaxies primarily involves understanding how baryons dissipated within the K cold dark matter (KCDM) paradigm (Freeman and Bland-Hawthorn, 2002). As a large fraction of baryons in galaxies are in stellar disks, an understanding of these disks is important. Over recent years, significant progress has been made by studying galaxy evolution as a function of space and time. This has been achieved through major extragalactic observational efforts (e.g., Tremonti et al., 2004; Kauffmann et al., 2003) and development of numerical simulations for large chunks of the universe (e.g., the Millennium simulation, Springel et al., 2005, 2006). These efforts have shown that the universe is a very dynamical place in which individual galaxies undergo accretion and merger events. In the late 1980s and early 1990s, the Milky Way was still largely seen as a rather static galaxy consisting of a few simple large building blocks: the disk, the bulge, and the halo. Since then, this picture has undergone dramatic changes. Although the change started earlier, the shift in view from a static to a dynamical Milky Way took on a new dimension when Ibata et al. (1994) pointed their telescope equipped with a multi-object spectrograph at star fields in the direction of the Galactic bulge. It was thought that the Bulge was essentially a spherical concentration of stars; the new observations were designed to further constrain the kinematics of Bulge stars and thus the Galactic potential in the inner regions. Ibata et al. (1994) found, in addition to the expected normal distribution of stellar velocities, an additional component with a completely different velocity. These stars turned out to be the brightest stars in a small galaxy in the direction of the stellar constellation Sagittarius, the Sagittarius dwarf spheroidal (dSph) galaxy. This small galaxy is currently being ripped apart by the Milky Way and its debris has now been traced across the sky (Belokurov et al., 2006). This and many similar studies, as well as direct imaging of nearby galaxies (e.g., Martínez-Delgado et al., 2010), provide direct evidence that late accretion occurs in regular galaxies such as the Milky Way. The Hipparcos satellite and its catalogue of stellar positions and proper motions for 118,218 stars (Perryman et al., 1997) led to the next steps in the full appreciation of the dynamical nature of the Milky Way. Several new streams of stars originating from dispersing star clusters and groups of stars with common space motions, probably shuffled together by the internal dynamics of our Galaxy, were found (Barrado y Navascues, 1998; Dehnen, 1998). It also became feasible to pre-select stars depending on their motion within the Milky Way for spectroscopic observations to obtain elemental abundances. Interestingly, spectroscopic studies of Hipparcos stars revealed that disk stars that are kinematically hot have different elemental abundances compared to stars on orbits similar to that of the sun. This seems to indicate different histories for the two stellar populations (e.g., Bensby et al., 2003, 2004; Reddy et al., 2006), but see Bovy et al. (2012) for a contradiction of this result. Apart from accretion events in the halo and complex disk structures, we have found that the central part of the Milky Way is not spherical. Indeed, the Galactic bulge hosts a bar (e.g., Stanek et al., 1994; Dwek et al., 1995; McWilliam and Zoccali, 2010) and complex stellar populations, perhaps even including an unexpectedly young stellar population (for a discussion see, e.g., van Loon et al., 2003; Bensby et al., 2013).

The history of a galaxy is imprinted in its stars, not only in their positions and velocities but also in their elemental abundances. The prospect of tracing the history of the Milky Way through the elemental abundances, ages, and kinematics of its stellar constituents is based on the fundamental assumption that the composition of elemental abundances in the outer atmosphere (the one we analyze) of low-mass stars does not change over time but instead forms a time capsule of the composition of the interstellar gas at the time the star formed (Lambert, 1989; Freeman and Bland-Hawthorn, 2002; Bland-Hawthorn et al., 2010). In addition, different elements are released to the interstellar medium by stars with different masses and therefore on different time scales. This means that the elemental abundance ratios measured in stellar spectra provide cosmic clocks that can be used to reconstruct the past history of star formation and gas accretion for the Milky Way. The classical clock is an a-element, such as oxygen or magnesium, relative to iron. Oxygen is produced in heavy stars and expelled into the interstellar medium when they explode as supernovae. This process acts on short time scales (Arnett, 1996). Iron is mainly produced in supernovae type Ia, which inject material into the interstellar medium on longer time scales (Nomoto et al., 1997). Thus, the ratio of these elements in the atmosphere of long-lived stars changes as a function of time. Elements mainly produced in asymptotic giant branch stars (AGB) enter the interstellar medium on yet another time scale; Simmerer et al. (2004) provided an example of how this is used to study the chemical evolution of the halo. 1.2. Differential analysis From the point of view of elemental abundances as tracers of the formation and evolution of the Milky Way, precision is more important than accuracy. This is fortunate, as systematic errors still can be substantial whereas it is feasible, with careful selection of targets, to obtain high precision results for elemental abundances in stars. Systematic errors can arise from inaccurate atomic data and simplifying assumptions in the modeling of stellar atmospheres and line formation. For example, stellar photospheres are assumed to be in local thermodynamic equilibrium (LTE), to be one-dimensional, and to have plane-parallel geometries. Current estimates shows that systematic errors are probably greater than 0.1 dex for solar-type stars (e.g., Asplund, 2005; Ruchti et al., 2013). However, for a homogeneous sample comprising stars with very similar stellar parameters, differential analysis can significantly increase the precision (Magain, 1984; Lambert, 1989; Gustafsson, 2004). Edvardsson et al. (1993), Nissen and Schuster (2010), and Meléndez et al. (2012) are examples of studies where this methodology is used, with increasing sophistication, in studies of F- and G-type dwarf stars leading to greater and greater precision. For example, Nissen and Schuster (2010) found internal errors as small as 0.04 dex, while Meléndez et al. (2012) achieved errors less than 0.01 dex. For B stars in the solar neighborhood, Nieva and Przybilla (2012) show internally very precise results. The current situation, whereby high precision is feasible but high accuracy is not, means that we should be cautious when comparing elemental abundances obtained for different types of object, such as dwarf and giant stars. We also need to be careful when

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comparing results from different studies of the same type of object, as methodological differences may lead to significant systematic differences. Even rather similar approaches can lead to differences that may obscure interesting structures in the abundance patterns. This review concentrates on our current understanding of elemental abundances in the stars in the Milky Way based on highresolution spectroscopy. We refer to Nissen (2013) for a brief introduction to abundance analysis of high-resolution spectra. We do not cover the origin of the elements; rather, we concentrate on the elements as tracers of the formation and evolution of the Milky Way. We focus mainly on the most recent studies that take advantage of advances in spectral coverage and resolution. It should be noted that we are careful when using the word ‘‘metallicity’’. We use it when talking about (i) the gross properties of stars or stellar populations, or (ii) a measure of the metal content of a star according to some spectral index or derived from photometry. This is denoted [M/H]. [Fe/H] then denotes the iron abundance for a star derived from Fe I and/or Fe II lines, where [Fe/ H] = log(NFe/NH)star  log(NFe/NH).

2. Characterizing the stellar disk(s) 2.1. In the solar neighborhood In their seminal study, Edvardsson et al. (1993) showed the power of combining elemental abundances in dwarf stars with kinematic information. They studied the elemental abundance trends in the stellar disk in a very systematic way. In particular, the stars had known kinematics and metallicity estimates based on Strömgren photometry. They selected approximately 25 stars for each metallicity bin (for defintion of their bins, see their Table 1). Since metal-poor stars have a lower space density in the solar vicinity, bins with the lowest metallicity are slightly less well populated. The concept of the thick stellar disk (Gilmore and Reid, 1983) was proposed before Edvardsson et al. (1993) published their study, but this was not taken into account for target selection. Key findings from this study include the tight abundance trends with spreads as small as 0.05 dex and a very flat trend for [Ni/Fe] as a function of [Fe/H] in the disk (Fig. 1). Owing to their target selection method, Edvardsson et al. (1993) observed only the a-enhanced part of the thick disk, confined to the lowest [Fe/H]. As a follow-up, Feltzing and Gustafsson (1998) studied elemental abundance trends for stars with super-solar iron abundances and

quantified the indications found by Edvardsson et al. (1993) for the most metal-rich stars in their sample. Feltzing and Gustafsson (1998) found that most elements showed a rather flat trend above solar metallicity but that, for example, [Na/Fe] notably increases with [Fe/H]. They found that this increase was not due to any particular orbit of stars, but seemed to be a common feature in the stellar disk. More recent studies of stars with super-solar metallicities by Pompéia et al. (2002a), Chen et al. (2003) and Trevisan et al. (2011) show similar results. Several studies took advantage of new kinematic data in the Hipparcos catalogue and explored the stellar disk in terms of populations. Studies of stars with kinematics typical of the thin and the thick disk (e.g., Chen et al., 2000; Reddy et al., 2003; Gratton et al., 2003; Bensby et al., 2003, 2005; Soubiran and Girard, 2005; Reddy et al., 2006; Reddy and Lambert, 2008) all revealed basically the same result: stars with kinematics typical of the thick disk differ in their elemental abundance ratios and ages from stars with kinematics typical of the thin disk. Fig. 2 gives two examples from Bensby et al. (2004) and Reddy et al. (2006). Fig. 3 shows an example of a different kinematic decomposition of the stars in the solar neighborhood (Gratton et al., 2003). Klaus Fuhrmann confirmed that this division between the two disks is real, at least in the immediate solar neighborhood. In a series of papers he explored the elemental abundances in a volumecomplete sample within 25 pc of the sun (Fuhrmann, 1998; Fuhrmann, 2000; Fuhrmann, 2004; Fuhrmann, 2008; Fuhrmann, 2011). The full sample contains stars further away. These mainly have lower metallicities and are generally associated with the thick disk or the halo rather than the thin disk. Fig. 4 shows the abundance trends for his volume-complete sample. Again, the stars, this time with no kinematic selection, split into two distinct abundance trends. The bottom panel in Fig. 4 shows the abundance trends but with magnesium as the reference element. Here it is perhaps even clearer that there are two different stellar populations present in the direct solar neighborhood. An a element is used as the reference instead of iron because iron has multiple stellar sources, while the a elements, and in particular oxygen, come from a single source (almost) and are therefore a better proxy for time (Wheeler et al., 1989). An obvious shortcoming of the various kinematic criteria used to define the thick and thin disks is that this approach assumes that the current orbit of a star reflects the orbital conditions of the cloud in which the star formed at the time of its formation. This is not necessarily a correct assumption. Sellwood and Binney (2002)

Fig. 1. Comparison of elemental abundance trends in the solar neighborhood as reported by Edvardsson et al. (1993) and Bensby et al. (in preparation). (a) As a function of [Fe/H] using data obtained by Edvardsson et al. (1993). (b) [Ti/Fe] as a function of [Fe/H] including only stars with 5600 < Teff < 6050 K obtained by Bensby et al. (in preparation). (c) [Ni/Fe] as a function of [Fe/H] using data obtained by Edvardsson et al. (1993). (d) [Ni/Fe] as a function of [Fe/H] including only stars with 5600 < Teff < 6050 K obtained by Bensby et al. (in preparation).

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0.4 0.2 0 -0.2 0.4 0.2 0 -0.2 0.4 0.2 0 -0.2 -1.5

-1

-0.5

0

[Fe/H]

Fig. 2. Left: Oxygen trends based on the [O I]630.0 line for stars investigated by Bensby et al. (2004) and Nissen et al. (2002). Data reported by Nissen et al. (2002) for thin and thick disk stars are denoted by open and filled triangles, respectively, and data for halo stars by asterisks. Data for thin and thick disk stars observed by Bensby et al. (2004) are marked by open and filled circles, respectively. Right: Data from Reddy et al. (2006) showing [a/Fe] as a function of [Fe/H] for thick disk stars (probability Pthick P 70%, top), thick/thin stars (50% < Pthick < 70%, middle), and thin disk stars (Pthin P 70%, bottom). Data for thin disk stars observed by Reddy et al. (2003) are also shown ().

Fig. 3. Elemental abundance trends for threep kinematically stellar compoffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffidefined ffi nents (Gratton et al., 2003). Thin disk, z2max þ 4  ecc2 < 0:35; dissipational component, VLSR > 180 km s1 and Rmax < 15 kpc. The accreted component comprises the remaining high-velocity stars.

Fig. 4. Data for stars within dSun < 25 pc based on analyses presented by Fuhrmann (1998, 2000, 2004, 2008, 2011). Final data courtesy of Klaus Fuhrmann. (a) [Mg/Fe] as a function of [Fe/H]. (b) [Fe/Mg] as a function of [Mg/H].

showed that owing to recurring interactions with spiral arms, it is possible for the orbit of a star to change such that it moves from one circular orbit to another circular orbit with a different galactocentric distance than that of the original orbit. This effect is called churning and its consequences for the kinematic properties of the stellar disk and on observed abundance trends can be significant (e.g., Schönrich and Binney, 2009b; Roškar et al., 2008). When the star moves to a new orbit it retains its energy and since the Galactic potential is weaker in the vertical direction (if the star moves to a larger RGal) the velocity vertical to the plane (WLSR) increases (but see Minchev et al., 2012). We would therefore expect the stars with the highest WLSR to be the oldest. Gratton et al. (2003) assigned population membership in a different manner from Bensby et al. (2003), Soubiran and Girard (2005), and Reddy et al. (2006). They used calculated stellar orbits to define three samples: a thin disk, a dissipative component, and an accreted component. The major reason for this division was that there is no

discernible difference between the elemental abundance trends in the halo and the thick disk at overlapping iron abundances (but see Section 4). Fig. 3 shows the elemental abundance trends for the three populations. If we apply the same selection criteria to a sample with greater numbers of thin disk stars, the thin disk trend becomes well populated. We also note that the dissipative component and the thin disk will overlap significantly in abundance space when using a sample that is well populated in the metal-rich part. Hence, it appears that this scheme perhaps gives most insight into the accreted component, which shows much greater scatter. Finally, Fig. 5 shows the eccentricity of the calculated orbits for a-enhanced and low-a stars in the volume-complete sample of Fuhrmann (1998, 2000, 2004, 2008, 2011). Although there are few a-enhanced stars, there is a clear difference between the two distributions. The different eccentricity distributions for stars with and without enhancement in a elements (Fuhrmann, 1998, 2000, 2004,

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Fig. 5. Cumulative distributions of eccentricities for stars with dSun < 25 pc from Fuhrmann (1998, 2000, 2004, 2008, 2011). The stars have been divided into two subsamples according to their [Mg/Fe] values, as indicated.

2008, 2011) suggest that in the thick disk, stars in the solar neighborhood have not been churned to their current positions, as they would then have had much more circular orbits. The same eccentricity pattern is observed for non-complete samples, such as data reported by Bensby et al. (in preparation) and Reddy et al. (2006). Note that this is not a selection effect: the stars were indeed selected to have non-circular orbits (i.e., lagging behind the solar rotation) but there was no knowledge regarding their a enhancement. They could all have turned out to have solar-like values for [a/Fe]. Bensby et al. (in preparation) searched for low-metallicity stars with typical thin-disk kinematics (below 0.7 dex) and found none. Stars with typical thin-disk kinematics typically have a lower [a/Fe] ratio than stars with typical thick-disk kinematics, also at the lowest metallicities.

2.2. Moving groups in the solar neighborhood and the concept of chemical tagging The idea behind the concept of moving groups is that stars form in clusters that later disperse as the clusters move within the gravitational potential of the Milky Way. These moving groups thus provide a potentially vital link between the open clusters and field stars. In a long series of papers, Olin Eggen identified a number of moving groups. He mainly defined moving groups as stars that share a common velocity in the direction of Galactic rotation, the V velocity. With the advent of the Hipparcos catalogue (ESA, 1997), much better kinematic information became available for stars in the solar neighborhood and a large number of studies explored the kinematic structure of the local disk. New moving groups or dissolving clusters were found (e.g., Barrado y Navascues, 1998; Arifyanto and Fuchs, 2006; Famaey et al., 2007; Klement, 2010; Bobylev et al., 2010). It has been confirmed that some moving groups are chemically homogeneous, while others are clearly dynamical structures in the stellar disk, probably induced by the bar. One moving group that has been confirmed is the HR 1614 group, first defined by Eggen (1992) and later confirmed, using Hipparcos parallaxes, by Feltzing and Holmberg (2000). High-resolution abundance results for HR 1614 obtained by De Silva et al. (2007) are reproduced in Fig. 6. This is clearly a

Fig. 6. Data for the chemically homogeneous moving group HR 1614 reported by De Silva et al. (2007). Differential elemental abundances relative to star HR 1614 as a function of the effective temperature (Teff) of the sample stars. Filled circles indicate stars identified as members based on their elemental abundances, while open symbols denote non-members. The other symbols indicate stars that were rejected as members based on their elemental abundances.

very homogeneous group of stars that share a common abundance pattern and hence a common origin. These findings have lent credibility to elaboration of the concept of chemical tagging (Freeman and Bland-Hawthorn, 2002; Bland-Hawthorn et al., 2010). Stars formed from the same cloud should share a common age and a common abundance pattern, representative of the cloud in which they formed. Thus, we should be able to use elemental abundances combined with age to pin down which field stars came from the same dispersing cloud. These ideas have led to the thought that the birth cluster of the sun could be identified if the right stars were selected. Brown et al. (2010) presented a first kinematic analysis based on Hipparcos data and metallicities based on data from Nordström et al. (2004). They provided a list of potential solar siblings, i.e. stars that were born together with the sun. It remains to be seen if these stars indeed are solar siblings (Cheng Liu and Gregory Ruchti, private communication). Not all moving groups are the result of a dissolved cluster. For example, the Galactic bar is able to shuffle stars in velocity space such that they appear to be a coherent group with a common origin (e.g., Dehnen, 2000; Antoja et al., 2011). Only a detailed analysis of the elemental abundances of the stars in such a group can reveal if they have a common origin. The Hercules stream (Famaey et al., 2005; Ecuvillon et al., 2007) was investigated by Bensby et al. (2007), who found that stars kinematically identified as belonging to this stream in fact show abundance patterns of both the thin and thick disk. In addition, the stars in the stream have a wide range of ages. The data indicate that this is probably a dynamical structure in the Milky Way disk and not a dissolving cluster of stars. Streams that belong to the halo may also appear in the solar neighborhood (Klement, 2010). One recently discovered example

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may cause offsets and extra scatter. Even if that were the case here, it is unlikely that the systematic differences would be sufficiently large such that the B stars really should be at the most metal-rich edge of the metallicity distribution function for F/G dwarf stars. The origin of the discrepancy in metallicities between the very youngest and moderately young stars from the thin disk in the solar neighborhood is not well understood. Potential explanations include the churning process discussed above (Sellwood and Binney, 2002). In the context of the following section, it is worth noting that the most metal-rich F/G dwarf stars in the solar neighborhood are on both circular orbits and orbits with large eccentricities, and hence their presence can probably not be ascribed to a single source. 2.4. Super-metal-rich stars: visitors from the Galactic bulge?

Fig. 7. Histograms showing the [Fe/H] distribution measured in the solar neighborhood for the low-Mg stars in Fig. 4 (i.e., data for F and G dwarf stars with dSun < 25 pc reported by Fuhrmann) (continuous line) and for nearby B stars (red shaded histogram) by Nieva and Przybilla (2012). The histograms have been normalized to a maximum height of 1 to emphasize the comparison of the spread in [Fe/H] and [Mg/H]. (For interpretation of the references to colour in this figure caption, the reader is referred to the web version of this article.)

is the Aquarius stream, which was identified in the RAdial Velocity Experiment (RAVE, http://www.rave-survey.aip.de/rave/) data set (Williams et al., 2011). This stream has a small [Fe/H] range and a narrow old age profile, both indicative of a remnant of a disrupted dwarf spheroidal galaxy that happens to pass through our neighborhood. Another example is the Arcturus moving group. Gilmore et al. (2002) found a sample of stars lagging behind the local standard of rest by 100 km s1 that they claim is associated with a disrupted satellite that merged with the Milky Way 10– 12 Gyr ago (Wyse et al., 2006). Navarro et al. (2004) suggest that these stars are the same group of stars that Eggen (1971) associated with the bright star Arcturus whose Galactic rotation also lags behind the local standard of rest by approximately 100 km s1. In a detailed abundance analysis, Williams et al. (2009) showed that the stars in the stream have a large range for the abundances for many elements, following the abundance trends seen in the disk in general. Therefore, this is not likely to be a moving group.

Stars with a wide range of metallicities exist in the solar neighborhood (Figs. 1–7). Stars with metallicities or iron abundances significantly greater than those of the sun have long attracted special attention (e.g., Spinrad and Luebke, 1970; Cayrel de Strobel, 1972; Grenon, 1972; Israelian and Meynet, 2008). The presence of such stars has implications for Galactic chemical evolution and our understanding of stellar evolution. Maeder (2002) provide a very brief overview of the multitude of issues related to metallicity. How metal-rich a star can be remains an open issue. It has even been argued that super-metal-rich stars may not exist, and there seem to be evidence that although dwarf stars in the local solar neighborhood have [Fe/H] ratios as high as 0.4–0.6 dex (compare Fig. 1 b and c and Fig. 8), the most metal-rich red giant stars may not exceed 0.2 dex (Taylor, 2001). It has been suggested that some of the super-metal-rich stars may be visitors from the Galactic bulge. Michel Grenon identified a sample of stars with high proper motions from the New Luyten Two Tenths (NLTT) catalogue for which stellar parameters were derived from Geneva photometry. Subsamples with high metallicities were then selected for high-resolution spectroscopy (Pompéia et al., 2002b; Trevisan et al., 2011). Pompéia et al. (2002b) found that the [O/Fe] trend followed that observed in previous studies of stars in the stellar disk in the solar neighborhood, in particular Nissen and Edvardsson (1992). Both Nissen and Edvardsson (1992) and Pompéia et al. (2002b) analyzed the forbidden [O I] line

2.3. How well mixed is the solar neighborhood? O and B stars provide very good tracers of recent elemental abundances in the gas from which the youngest stellar disk stars formed. Recent studies of such stars in the direct solar neighborhood reveal very tight distribution functions for a range of elements (Nieva and Przybilla, 2012). Fig. 7 shows their Fe and Mg distributions compared to data for low-a stars reported by Fuhrmann. The young stellar sample has a very tight distribution for Fe and Mg abundances compared to the much wider range for low-a stars. If we consider that these young stars are the result of recent star formation from enriched gas in the direct solar neighborhood, then it is clear that the solar neighborhood also contains older stars that are more enriched in iron than these young stars. Hence, the recent star formation point to great homogeneity in the gas from which the young stars formed, while the somewhat older stars point to a much more complex star formation history. We cautioned in the Introduction that it is dangerous to compare results from different types of stars as systematic differences

Fig. 8. Elemental abundances reported by Trevisan et al. (2011) for stars in the solar neighborhood selected as being metal-rich or super-metal-rich based on Geneva photometry and elemental abundances for micro-lensed stars in the Galactic bulge from Bensby et al. (2013). (a) [Si/Fe] and (b) [Ni/Fe] as a function of [Fe/H]. Black dots denote data from Trevisan et al. (2011) and red squares denote data from Bensby et al. (2013). (For interpretation of the references to colour in this figure caption, the reader is referred to the web version of this article.)

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Fig. 9. [Si/Fe] as a function of VLSR, Bensby et al. (in preparation) and Trevisan et al. (2011). a) Data for all stars reported by Bensby et al. (in preparation). b) Data for all stars reported by Trevisan et al. (2011). All stars reported by Bensby et al. (in preparation) with c) [Fe/H] > 0 and d) [Fe/H] > +0.25.

at 630.0 nm, which is one of the few lines that does not suffer from non-local thermodynamic equilibrium (NLTE) effects (Kiselman, 1991; Kiselman, 2001). Interestingly, Trevisan et al. (2011) showed that all the stars in their sample have solar levels of a enhancement (compare Fig. 8). Fig. 9a essentially reproduces Fig. 23 of Trevisan et al. (2011) but instead of data from several studies we use the stars reported by Bensby et al. (in preparation). This shows that there are stars with large and small [Si/Fe] ratios at all low VLSR and that there is a gap in the elemental abundances such that there is one group with low and one with high [Si/Fe]. The subsequent panels show the same plot for all stars reported by Trevisan et al. (2011) and for stars reported by Bensby et al. (in preparation) with [Fe/H] > 0 and +0.25 dex, respectively. Although it would be possible to conclude from the data of Trevisan et al. (2011) that stars with super-solar metallicities also have a low-VLSR kinematic signature, data from Bensby et al. (in preparation) reveal that stars with super-solar metallicities have a range of VLSR values and the signature is a consequence of biases in the selection of targets. However, this does not exclude the possibility that some of the super-metal-rich stars in the solar neighborhood originate from the inner disk, and perhaps even the Galactic bulge. The stars selected for Fig. 9d) have a mean age of 3.2 Gyr with scatter of 1.3 Gyr, which is an age typical of the younger thin disk in the solar neighborhood (only stars with absolute errors in age less than ±2 Gyr are included in this calculation). Among all stars more metal-rich than the sun in the sample selected by Bensby et al. (in preparation) there is just a handful with age greater than 5 Gyr. These do not have a specific kinematic signature but are evenly spread over the full range of VLSR values seen for the low-[Si/Fe] stars. We can conclude that metal-rich and super-metal-rich stars present in the solar neighborhood are not distinct in terms of their kinematics or their ages, but seem to form a continuous distribution. 2.5. Beyond the immediate surroundings So far we have essentially considered stars in the direct solar neighborhood. Given that most stars studied to date are within

approximately 100 pc of the sun, the stellar disk extends to a few kpc above the Galactic plane and we are 8 kpc from the Galactic centre, any conclusions drawn from these data must await confirmation based on stars that are either in a different part of the Galaxy (in situ measurements) or are known to have travelled here from their formation site. Gradients in elemental abundances are an efficient and interesting constraint on chemical evolution models (see for example discussions in Lambert, 1989; Chiappini et al., 1997; Cescutti et al., 2007). Such gradients can be traced in several ways, all of which impose their own constraints on models. While studies of abundances in H II regions constrain the gradients in the gas today, studies of Cepheids and old dwarf and giant stars (both in the field and in open clusters) probe the gradients as a function of time. High-resolution spectroscopy of faint stars is very time-consuming and requires large telescopes. Therefore, most high-resolution studies have focused on nearby stars. However, current telescopes are capable of obtaining high-resolution spectra of stars as faint as V = 16–17 with a decent signal-to-noise ratio within an exposure time of a few hours. Stellar spectroscopists have not extended their studies from field stars in the solar neighborhood to field stars elsewhere for two main reasons: lack of distance information and lack of large-scale photometric databases from which stars can be reliably selected. The latter has greatly improved with ongoing VISTA surveys, for example, which provide high-quality photometry for the whole Southern sky and crucially for the disk. Reliable distances based on parallaxes will come with Gaia, but there are other means to find suitable objects. For example, the VISTA database has been used extensively to select turn-off targets studied in the Gaia-ESO Survey (Gilmore et al., 2012) out to several kpc. For greater distances the main focus is still, however, on open clusters and specific types of stars that are easily identifiable and have a fairly narrow range of absolute magnitude or other means to derive their distances (e.g., red clump stars and Cepheids). 2.5.1. Old and young stellar components For open and globular clusters it is relatively easy to derive a reliable distance estimate based on photometry. This makes open clusters attractive targets for studying the stellar disk at greater distances. Approximately half of the know open clusters have reliable distance estimates (Dias et al., 2002; Dias et al., 2012). It is also possible to derive ages for the open clusters from color-magnitude diagrams, so temporal changes in abundance gradients can be investigated (e.g., Yong et al., 2012). Analysis of high-quality

Fig. 10. [Fe/H] for open clusters normalized to solar [Fe/H] as a function of the galactocentric radius measured by different authors (taken from Lépine et al., 2011). Only results based on high-resolution spectroscopy are included. [Fe/H] errors were taken from the original references. Galactocentric distances were recalculated for RSun = 7.5 kpc (the short distance scale favored by Lépine et al., 2011). Different colors correspond to the following age ranges: blue, <200 Myr; green, 200–1200 Myr; red, >1200 Myr. The dashed horizontal line indicate the average metallicity on both sides of corotation.

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Fig. 11. [Fe/H] for Cepheids as a function of the galactocentric radius (Taken from Andrievsky et al. (2013) who used data from Luck and Lambert (2011)).

Fig. 12. [a/Fe] for Cepheids as a function of the galactocentric distance, taken from Lépine et al. (2011). The galactocentric distance is on the short distance scale favored by Lépine et al. (2011), with RSun = 7.5 kpc.

spectra for a few giant stars in each cluster has shown that there is probably a rather steep [Fe/H] gradient that flattens at greater galactocentric distance (examples of studies finding this gradient and flatting include Lépine et al., 2011; Yong et al., 2012 see also Fig. 10). For Cepheids such a plateau is perhaps less discernible (cf. Fig. 11), although the Cepheid data are less complete at large galactocentric distances (Luck and Lambert, 2011; Yong et al., 2012); however, Lépine et al. (2011) found a difference for Cepheids at the corotation radius. In the outer disk the open clusters show a mean [Fe/H] of approximately –0.3 to –0.4 dex (Fig. 10), but their a elements are enhanced in relation to iron, albeit with large scatter (Yong et al., 2012). Fig. 12 shows [a/Fe] for a large compiled sample of Cepheids (Lépine et al., 2011). These stars show a very flat and rather tight abundance trend as a function of galactocentric distance, which ap-

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pears to be in notable contrast to observations for the open cluster in the outer disk (which essentially are all old). Thus, this first seems to indicate that the star formation history in the outer part of the disk has changed over time. A canonical interpretation would be that the open clusters represent a fast star formation period, whereas the Cepheids, which formed later, are part of a more sedate period of star formation that does not reach high [a/Fe] ratios (see, e.g., McWilliam, 1997; Matteucci, 2001). However, it is not certain that this discrepancy is real. The open clusters in the outer disk are rather metal-rich (Fig. 11) and in almost all cases only very evolved red giant stars have been studied in the high-resolution abundance studies discussed here. The analysis of sepctra for evolved, metal-rich stars is problematic as their spectra are very crowded by molecular lines, adding to the multitude of atomic lines (e.g., Fulbright et al., 2006 and see the discussion in Section 3). A recent analysis of stellar spectra for stars at different evolutionary stages belonging to the open cluster M 67 lends further support to a cautious interpretation of the high [a/Fe] ratios found in open clusters at large galactocentric distances. Önehag et al. (2011) analyzed a solar twin in M 67, i.e. a dwarf star with stellar parameters very similar to those of the sun, and found an [Fe/H] of +0.02 dex and [X/Fe], where X is several elements including a elements, to be within 0.03 dex of the solar values. This is in contrast to analysis of evolved giant stars in the same cluster for which Yong et al. (2005) found an [Fe/H] very close to the solar value but with abundance ratios that differ significantly from the solar values. The number of studies of (likely) old field stars in the outer disk are few. Perhaps the most interesting of these compares elemental abundances trends for (warm) red giant stars in the inner disk, solar neighbourhood and the outer disk. The results are reproduced in Fig. 13. Before discussing these results in the context of open clusters, it is important to note that outer disk stars do not reach quite as far out in the disk as open clusters and that they are situated somewhat above the mid-plane of the Milky Way disk. Nevertheless, it is still very interesting to note that stars chosen only on the basis of their position in color-magnitude diagrams (i.e. evolutionary stage) and their distance show a remarkably tight abundance range for both [Fe/H] and [a/Fe] in the outer disk. The small spread in iron abundance is very reminiscent of what we see for the open clusters (Fig. 11 and Yong et al., 2012). The lack of high-a stars appears to be real, as an analysis of an identical sample in the inner disk (Fig. 13 left) shows elevated a abundances. Preliminary results from the Gaia-ESO Survey appear to confirm these observations.

Fig. 13. Elemental abundance trends for K giant stars (Bensby et al., 2010, 2011; Alves-Brito et al., 2010) as a function of galactocentric radius (as indicated on top of the columns). Left: stars in the inner disk. Middle: stars in the thin disk () and thick disk () in the solar neighborhood as defined by Alves-Brito et al. (2010). Right panel: stars in the outer disk. The lines are the same in all three panels and are drawn by hand to guide the eye (Bensby et al., 2011). Figure courtesy of Thomas Bensby.

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The gradients at earlier times can be studied by looking at young A and B stars. These heavier stars are intrinsically bright so they can be studied at large distances. Recent work by Norbert Przybilla and Maria Fernanda-Nieva has shown some remarkably tight abundance trends as a function of galactocentric distance. To the best of our knowledge this work has not been published, but impressive results have been presented at several conferences.

2.5.2. Radial abundance gradients in other galaxies Radial abundance gradients are not confined to the Milky Way and have been found in several other galaxies. For very nearby galaxies, such as Andromeda, it may still be feasible to obtain abundance data from stellar clusters or young bright stars. Abundance measurements for the stellar component of more distant galaxies is more difficult. However, abundances in gas are readily observable. Bresolin et al. (2012) measured radial gradients in H II regions in four nearby disk galaxies (M 83, NGC 1512, NGC 3621, and NGC 4625). All show a steeper gradient first and then a flat abundance trend to the largest radii (Fig. 11 in Bresolin et al. (2012) summarizes their results). The qualitative shape of these trends is very similar to that observed for open clusters and Cepheids in the Milky Way (Lépine et al., 2011). Jones et al. (2013) investigated the time evolution of radial metallicity gradients in disk galaxies by observing a set of lensed spiral galaxies at z = 2  2.14. Comparison with local data for other galaxies as well as the Milky Way showed that the abundance gradients probably evolve over time within extended disk galaxies such that the gradient is much steeper at earlier times and then flattens at later times as the radius of the galaxy increases (Jones et al., 2013).

2.5.3. Above the Galactic plane To the best of our knowledge, there is no high-resolution abundance study of field dwarf stars situated well above the Galactic plane. Large-scale surveys such as GALAH and the Gaia-ESO Survey will change this situation. For example, the Gaia-ESO Survey will provide more than 1000 dwarf stars observed with VLT/UVES (high resolution) within 1 kpc of the sun. Such a sample does not currently exist, and it is likely that unexpected results will be obtained.

2.6. Formation scenarios: a brief overview The observations reviewed so far point to a rather complex structure for the Milky Way stellar disk. In addition to the work discussed above, several studies of large sets of photometric data and gross abundance data derived from low-resolution spectra have shaped our current knowledge of the chemodynamical status of the Milky Way. Interesting examples include studies by Bovy et al. (2012), Liu and van de Ven (2012), and Schlesinger et al. (2012). As discussed by Freeman and Bland-Hawthorn (2002), for example, it is our hope that data such as those reviewed here eventually will help us to decipher the formation and evolution of the Milky Way. Are we there yet? We use the formation of the thick disk as an example of the current ability of observations to constrain models and of models to explain observations. There are currently at least six different scenarios available in the literature to explain the origin of the Milky Way thick disk. All scenarios have been subjected to modeling efforts. Here we briefly review the most salient points of the models and the ability of the models to explain all or part of the data and how the models can be challenged by the data.

2.6.1. Monolithic dissipative collapse In this scenario the stars that today occupy the thick disk were formed during dissipative collapse of a proto-galaxy subsequent to the formation of the stellar halo (e.g., Eggen et al., 1962; Larson, 1974). At least one detailed simulation has been carried out for this scenario (Burkert et al., 1992). Although the collapsing gas disk still had a large vertical size, the high star formation rate and/or associated energy feedback by supernova explosions prevented further collapse of a disk, leaving thick disk stars with a large velocity dispersion and large scale height. This scenario can explain the old age and high [a/Fe] abundance of the thick disk, whereas a dissipative process itself may leave a finite metallicity gradient in the stellar system, contrary to suggestions from currently available observational results. 2.6.2. Multiple dissipative mergers at high redshift In line with the current standard scenario of structure formation (e.g., Springel et al., 2005), the thick disk stars formed from multiple mergers of gas-rich building blocks at high redshift, which eventually leaves an old stellar disk with high velocity dispersion. Like the monolithic collapse model, this scenario can explain the old age and high [a/Fe] abundance of the thick disk (Brook et al., 2004, 2005). Further studies are needed to prove that a chaotic process involved in multiple dissipative mergers reproduces the systematic decrease reported for the disk’s mean rotational speed with height above the plane. 2.6.3. Tidal debris of shredded satellites A significant part of the thick disk could have originated directly from the tidal debris of shredded satellites when they accreted onto a forming galaxy through decay and circularization of the orbits by dynamical friction (Abadi et al., 2003). The fraction of the tidal debris is greater for the older (>10 Gyr) thick-disk component. This scenario can reproduce the dynamical structure and age distribution of the thick disk. However, it is yet uncertain how thick disk stars with high metallicity, high [a/Fe], and old age have formed in such shredded satellites, in contrast to the properties of stars in the current Galactic satellites actually observed (Tolstoy et al., 2009). 2.6.4. Disk heating triggered by minor merger In this scenario a pre-existing stellar disk in the form of the thin disk was dynamically heated during a minor merger with a luminous satellite galaxy, leaving a stellar disk with large velocity dispersion and large scale height (e.g., Quinn et al., 1993; Velazquez and White, 1999; Villalobos and Helmi, 2008). This scenario predicts no metallicity gradient due to dynamical mixing of stars and can also explain the presence of the vertical gradient observed for the mean rotational speed. It is also noted that debris stars from a tidally disrupting satellite galaxy themselves make an additional contribution to thick disk stars in the form of a bunch of eccentric orbits. However, prediction of the latter might be in disagreement with the orbits observed for disk stars (Sales et al., 2009), although current data sets are too small and errors too large for firm conclusions, so future large-scale surveys and Gaia proper motion data are necessary to settle this question. The elemental abundance trends in the resulting thick disk would probably differ from those of the current thin disk and the stars would be older. Hence, it appears likely that the age and elemental abundances observed in the current thin and thick disk in the solar neighborhood could be explained in this scenario. To the best of our knowledge, no model with sufficient detail for detailed comparisons with observations of elemental abundances is available. Disk heating is also driven by minor merging of CDM sub-halos or dark satellites (Hayashi and Chiba, 2006; Kazantzidis et al., 2009). The standard CDM theory suggests the presence of a large

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such that meaningful comparisons can be made. In particular, many of the observational studies discussed in the previous subsections are still limited to relatively small number of targets. Large-scale studies of stars with high quality data are necessary to improve the comparisons further; Section 5 lists the most important future surveys.

Fig. 14. An example of growth in disk thickness (from left to right) driven by a merger of CDM sub-halos (Hayashi and Chiba, 2006).

number of sub-halos in a galaxy-sized halo (Klypin et al., 1999; Moore et al., 1999) that can interact with a luminous stellar disk and cause dynamical heating, leaving a hot disk system similar to the thick disk; Fig. 14 shows an example. In this scenario, since dark satellites are involved in disk heating, the effect of tidal debris of shredded satellites in the chemokinematics of the thick disk formed, especially the presence of stars with high orbital eccentricities and low [a/Fe] ratios, can be ruled out. Further studies are needed on the timing of the merger process for reproducing the metallicity and age distributions observed for stars. 2.6.5. Radial migration of disk stars Without relying on accretion from outside, the presence of thickdisk-like stars in the solar neighborhood can potentially be explained by an internal dynamical process in the Galactic disk. Radial migration of disk stars is expected to occur through scattering by transient spiral arms and/or molecular clouds (Sellwood and Binney, 2002). Thus, stars in the solar neighborhood could be a mixture of those formed nearby and those transferred from inner/outer disk regions. This scenario has been extensively studied in dynamical models (e.g., Roškar et al., 2008). The consequences of radial migration for the abundance distributions observed in the solar neighborhood have been explored (e.g., Schönrich and Binney, 2009a,b). Although the general idea still holds, there is no consensus on whether thickening of the disk arose from this churning process or if it had a different origin; for example, Minchev et al. (2012) showed that radial migration is not necessarily accompanied by thickening of the disk. No doubt future studies such as the Gaia-ESO Survey (Gilmore et al., 2012) will lead to progress in this field. 2.6.6. Clumpy disk formation at high redshift Simulations of gas-rich young galaxies focus on the formation of internal clumps by gravitational instabilities (e.g., Noguchi, 1999; Bournaud et al., 2007). These clumps give rise to disk thickening via strong stellar scattering. The resulting internally heated disk may correspond to the thick disk observed in the Milky Way and many other galaxies (Bournaud et al., 2009). A thick disk formed via such internal processes has a constant scale height with galactocentric radius (Bournaud et al., 2009) and is expected to show a symmetric vertical distribution of stars with respect to the disk plane. Further studies of these disk properties are needed to test this model. Work by Bovy et al. (2012), for example, may indicate that scale height is a function of stellar properties including elemental abundances, although Liu and van de Ven (2012) obtained a different result for the same data. Although it is likely that all of the processes discussed here occur to some extent in the Milky Way, it is not clear which, if any, is the dominant mechanism. Comparison of model predictions and observations is challenging for two reasons: (1) the models sample stellar populations at a coarse scale owing to low resolution in the simulations; and (2) the stellar samples are not always defined

2.6.7. Galactic chemical evolution So far we have only briefly mentioned Galactic chemical evolution. The aim of the abundance studies we reviewed is to constrain the formation and evolution of the Galaxy, particularly Galactic chemical evolution. This requires comparisons with models. A thorough review of models of Galactic chemical evolution is beyond the scope of our article. The framework for Galactic chemical evolution is well established and has been reviewed many times; it is described in detail in the book by Francesca Matteucci (Matteucci, 2001). Models of Galactic chemical evolution have successfully explained the origin of the overall elemental abundance trends in different populations and metallicity distributions in the solar neighborhood, and have shown the need for processes such as infall (e.g., Wheeler et al., 1989; Timmes et al., 1995; Chiappini et al., 1997). However, as reviewed here, the observational evidence is increasingly rich. The combination of elemental abundances, stellar ages and kinematic data and increasing computing power is leading to a new era for explorations in which all of these can be combined in models. Notable examples of chemodynamical modeling include a study by Kobayashi and Nakasato (2011), who used a self-consistent hydrodynamic code including supernova feedback and chemical enrichment. This code can, for examle, be used to explore elemental abundance trends for different stellar components, such as the solar neighborhood and the Galactic bulge. A different approach was developed by Minchev et al. (2012), who combined a classical chemical model with a rendition of the formation of a galaxy in KCDM to explore the formation of, for example, the thick disk via mergers and radial migration. In a series of papers Chris Brook and collaborators (Brook et al., 2012, and references therein) added chemistry to their particles to follow the buildup of iron and a elements within the different resulting populations. In spite of our enthusiasm for the recent promising developments in chemodynamical modeling we feel that simple explorations still have an important role to play and they can, yield new fundamental insights. A recent illustration of this is the investigation of potential origins of metallicity gradients in the Galactic bulge by Martinez-Valpuesta and Gerhard (2013). Although this is not a chemical evolution model, it nicely illustrates that simple studies can give further insights. Martinez-Valpuesta and Gerhard (2013) showed that if the stellar disk in the inner Galaxy prior to bar formation had a radial metallicity gradient, the dynamical evolution that occurred when the bar formed naturally would lead to metallicity gradients in the Galactic bulge (as projected onto the observed plane) similar to those found in analyses of VVV and 2MASS data for the Bulge (Gonzalez et al., 2013). Thus, the gradients observed are not necessarily proof of a monolithic collapse. Stellar yields, the return of elements to the interstellar medium in the form of winds from heavy stars and supernova explosions, are an essential ingredient in any model of Galactic chemical evolution (Arnett, 1996). It is thus perturbing that there appears to have been little recent progress in the certainty in the calculations of yields from supernovae and heavy stars (Romano et al., 2010). 3. Elemental abundance trends in the Galactic bulge Bulges in spiral galaxies were originally thought to be rather spherical objects and, since they have red colors, should contain

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mainly old stars. The Galactic bulge was envisioned to be no different. Indeed, most observations indicated a metal-rich and old stellar population (e.g., Rich, 1988; Terndrup, 1988). However, over the last decades the picture of bulges in other galaxies and in our own Galaxy has become more complex. In particular, we are now certain that the Milky Way contains a substantial bar (Stanek et al., 1994; Dwek et al., 1995). The bar may in fact make up the largest fraction of the Bulge, with little room for a classical, spherical bulge (Shen et al., 2010). Rich (2013) presented an extensive review of past and present views on the Galactic bulge based on all types of observations. Studies of the solar neighborhood have shown that it is reasonably straightforward to derive precise elemental abundances from spectra of dwarf stars and these can be used to trace the chemical history of Galactic populations. However, the Galactic bulge is situated sufficiently far away that Bulge dwarf stars at the turn-off have V  19–20 (Holtzman et al., 1993; Ortolani et al., 1995; Feltzing and Gilmore, 2000). This is too faint to obtain high-resolution spectra, even with 8- and 10-m-class telescopes. Studies of intrinsically bright giants showed that the Bulge is metal-rich and enhanced in a elements, indicative of a short formation time. (important examples of Bulge studies include, McWilliam and Rich (1994), Fulbright et al. (2007), Lecureur et al. (2007)). In these studies the high [a/Fe] continued to super-solar metallicities, setting the Bulge apart from the solar neighborhood, which has solar abundances at these metallicities (Bensby et al., 2005). However, the spectra of giants are notoriously difficult to model owing to the presence of large numbers of molecular lines. Past studies of metal-rich stars in the Galactic bulge and the metal-rich globular clusters NGC 6528 and NGC 6553 illustrate this point. For the globular cluster NGC 6553, Barbuy et al. (1999) obtained [Fe/H] = 0.55 dex for two stars with Teff = 4000 K, while Cohen et al. (1999) found [Fe/H] = 0.16 dex and enhanced [a/Fe] for five horizontal branch stars that are warmer (Teff = 4700–4800 K). Cohen et al. (1999) concluded that a difference of 0.1–0.15 dex can be accounted for by the different stellar atmosphere models used and differing atomic line data, but they also found that the remaining discrepancy is most likely due to the different (and more severe) degree of molecular absorption in the coolest stars. Furthermore, Barbuy (1999,) found that some, but not all, a elements were very elevated relative to iron in these two clusters. In studies of metal-rich K giants in Baade’s Window by Rich (1988), McWilliam and Rich (1994), and Fulbright et al. (2006), the metallicities derived changed because of increasing resolution of the observed spectra, and subsequent inclusion of molecules in the spectral analysis. Alves-Brito et al. (2010) performed the first large differential study for giant stars in the Bulge and in the solar neighborhood and found very similar results for abundance trends and that a elements are not enhanced for stars with super-solar metallicities. Fig. 15 shows their results for magnesium, from which it is clear that the thin disk in the solar neighborhood differs from the Bulge population. A more extensive comparison would be possible if we could observe dwarf stars in the Galactic bulge. This is feasible if we take advantage of micro-lensing events, during which a star can be magnified by several magnitudes (Mao, 2012), making it possible to obtain a high-quality spectrum in a few hours. Johnson et al. (2007, 2008), and Cohen et al. (2008) obtained the first high signal-to-noise spectra of micro-lensed dwarf stars in the Galactic bulge. The results were surprising: the stars had very high metallicities and solar abundances for a elements. Bensby et al. (2013) have executed a Target-of-Opportunity program on the VLT and took advantage of additional observation possibilities on the Magellan and Keck telescopes. Results from these studies are shown in Figs. 8, 15, and 16. Overall, the abundance trends for Bulge dwarfs are very tight and the Bulge stars follow the upper

Fig. 15. Top: Plot based on data presented by Alves-Brito et al. (2010) showing elemental abundances for the Bulge (black dots) and thin (blue dots) and thick (red circles) disk giants. Bottom: Elemental abundances for thin (blue dots) and thick disk (red circles) dwarf stars (Bensby et al., in preparation) and for micro-lensed Bulge dwarf stars (Bensby et al., 2013). (For interpretation of the references to colour in this figure caption, the reader is referred to the web version of this article.)

Fig. 16. [O/Fe] as a function of [Fe/H] for micro-lensed dwarf stars in the Galactic bulge (Bensby et al., 2013). The symbol size indicates the age of the stars.

envelop of the elemental abundance trends traced by local, kinematically selected, thick disk stars. These results are for the inner regions of the Galactic bulge. Several studies have targeted fields further out, exploring possible trends along minor and major axes. As summarized by Johnson et al. (2013), all studies agree that abundance trends for different positions in the Bulge show the same overall elevated [a/Fe] levels below approximately –0.3 dex in [Fe/H], with a declining trend towards higher metallicity. Elemental abundances should preferably be combined with stellar ages and kinematics to better constrain the current properties of the Galactic bulge and its chemical history. Deep color-magnitude diagrams of the Bulge population show a faint and red turnoff, compatible with a metal-rich and uniquely old population (e.g., Feltzing and Gilmore, 2000; Clarkson et al., 2008). By contrast, studies of micro-lensed dwarf stars in the Bulge also revealed some younger ages (Bensby et al., 2013, and Fig. 16). Bimodal age distributions might be possible if, for example, bars have the ability to restart star formation in a bulge. Coelho and Gadotti (2011) found that bulges in spiral galaxies with bars show stellar populations with bimodal age distributions, while galaxies without bars do not. It is worth noting that in chemodynamical simulations it is possible for Milky Way-like galaxies to sustain a low level of star formation over extended times (Kobayashi and Nakasato, 2011).

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Chemical evolution models can help in making sense of the rather complex situation in the Galactic bulge. Interesting studies include Cescutti and Matteucci (2011) that showed how the initial mass function can be constrained using the observed metallicity distribution function for the Bulge. In particular, they found that a top-heavy initial mass function was necessary to understand the then current metallicity distribution function for giant stars in the Galactic bulge. Further studies of the metallicity distribution functions such as that by Ness et al. (2013) may change this picture. Another interesting study is provided by Rahimi et al. (2010) who showed that by combining elemental abundances and kinematics with numerical studies it is possible to disentangle different formation scenarios for the Bulge (e.g., secular vs. merger origin). Infrared observations offer the possibility to investigate stars in highly reddened regions. Origlia and Rich (2004) and Ryde et al. (2010) provide examples of studies taking advantage of infrared spectroscopy in the Bulge. More recently, Rich et al. (2012) confirmed the high a values also for the innermost regions of the Bulge (l 6 1° and b = 0°, 2°.65). APOGEE will deliver the next major step in infrared spectroscopy (Allende Prieto et al., 2008). Future largescale optical and infrared studies of the Bulge will include 4MOST and MOONS (de Jong et al., 2012; Cirasuolo et al., 2012). In the meantime, the upgraded CRIRES spectrograph (Oliva et al., 2012) will facilitate targeted deep studies in the inner Bulge.

4. Characterizing the stellar halo 4.1. Multiple halos: a long history of tracing the evidence Stars in the Galactic halo were first identified as stars with high velocities and an asymmetric velocity distribution in the azimuthal direction of the disk in the solar neighborhood (Oort, 1922). Most of these stars were later found to have large UV excess compared to stars in the Hyades cluster. Large UV excess is indicative of weak metal lines (low metallicity) (Roman, 1955). The connection between stellar kinematics and elemental abundances was first analyzed in the context of the formation of the Galaxy in a seminal paper by Eggen et al. (1962). The stellar populations in the Galactic halo, including globular clusters, are now characterized as mainly metal-poor ([Fe/H] < 1) and dominated by random motions, in contrast to stars in the disk, which show an orderly motion. It has been realized since the 1980s that the halo component actually shows multiplicity in the properties of its stars. Searle and Zinn (1978) showed that globular clusters in the halo have a finite diversity in their horizontal-branch morphology, which was interpreted to indicate the presence of relatively young globular clusters in the outer halo (Zinn, 1993). This finding, as well as the lack of a metallicity gradient among halo clusters, led to the idea that the galaxy formation process was not a simply monolithic collapse but was accompanied by accretion of small building blocks, as envisaged by Searle (1977). Metal-weak field stars in the halo also appear to have a nonmonotonous spatial structure. The velocity distribution for halo stars in the solar neighborhood shows radially elongated velocity dispersion, representing an anisotropic velocity field. According to dynamical models, this suggests that the shape of the stellar halo is highly flattened in the radial direction (e.g., White, 1985; White, 1989; Levison and Richstone, 1986; Ratnatunga and Freeman, 1985). However, direct star-count analyses suggested that the stellar halo is spherical (Bahcall and Soneira, 1984; Wyse and Gilmore, 1989). Which picture is the correct one? Both would actually be correct if the halo has a dual structure: Hartwick (1987) showed that the metal-poor RR Lyrae stars delineate a two-component halo, with a flattened inner component and a spherical outer

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component. Sommer-Larsen and Christensen (1989) also proposed, based on their analysis of radial velocities for distant blue-horizontal branch (BHB) stars, that the stellar halo is better represented by the presence of a spherical outer component in addition to a flattened inner halo. Further evidence of multiplicity in the halo has emerged in recent studies of halo samples near the Sun. Using data for 118 nonkinematically selected metal-poor stars, Sommer-Larsen and Zhen (1990) constructed a global density distribution for the halo based on superposition of sample orbits. They found that the halo consists of two components: a main, nearly spherical component with an axial ratio of approximately 0.85 comprising the large majority of the mass in the stellar halo; and a highly flattened component, contributing approximately 40% of the density of the halo at the position of the Sun. Chiba and Beers (2000) carried out the same orbital analysis for a much larger (1200) local sample of nonkinematically selected stars using a combination of Hipparcos data and several ground-based proper motion studies. They arrived at the same two-component picture for the halo structure. More recent large surveys such as the Sloan Digital Sky Survey (SDSS, http://www.sdss.org) offered important opportunities to revisit this subject. Carollo et al. (2007), Carollo et al. (2010) analyzed the kinematics of SDSS calibration stars and found that a broadly overlapping, two-component halo is needed to explain their chemokinematic properties: a flattened inner halo with zero or some pro-grade rotation; and an outer, more spherical and more metal-poor halo with retrograde rotation (see also Schönrich et al., 2011; Beers et al., 2012). The counter-rotating nature of the outer halo is supported by several recent kinematic studies of distant BHB and RR Lyrae stars (e.g., Deason et al., 2011a; Kinman et al., 2012; Kafle et al., 2013; Hattori et al., 2013). The assembly and analysis of these halo tracers provide a global halo density distribution that shows a break in slope at approximately r = 30 kpc, representing a two-component halo structure (Watkins et al., 2009; Sesar et al., 2011; Deason et al., 2011b). Additional evidence of the multiplicity of the halo has been provided by An et al. (2013). Using a metallicity calibration from SDSS photometry, they found that candidate halo stars with retrograde motion are more metalpoor than those with pro-grade motion, which indicates that the outer halo is more metal-poor than the inner halo.

4.2. Insight from high-resolution abundance studies The probable double nature of the halo should also be evident in detailed elemental abundance data for stars, as we expect the enrichment of elements to vary according to the formation process for each halo component, reflected in observable differences in the chemodynamical structure. More specifically, the elemental abundances for stars in the metal-poor, counter-rotating outer halo may show systematic differences to those in the inner, pro-grade halo. Several studies of elemental abundances in halo stars have indeed found evidence that not all stars that belong to the halo show the same canonical high values for [a/Fe], for example. Stephens and Boesgaard (2002) reported that stars with large apo-centric distances tend to have lower [a/Fe] ratios, while no correlations between abundances and other orbital parameters were identified. Fulbright (2002) found the lowest [Mg/Fe] ratios for stars with the largest rest-frame velocities. It is worth noting that none of the stars could be easily associated with the Sagittarius dSph galaxy or other potential streams resulting from either globular clusters or accreted dwarf galaxies. These studies of the elemental abundances of nearby stars taking into account their kinematic properties provide important insights into the formation of the Milky Way halo (Gratton et al., 2003; Roederer, 2009; Zhang et al., 2009; Ruchti et al., 2010; Ishigaki et al., 2010; Ishigaki et al., 2012).

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Fig. 17. [Mg/Fe] and [a/Fe] versus [Fe/H] as reported by Nissen and Schuster (2010). Crosses and circles/triangles denote data for thick disk and halo stars, respectively. Halo stars with data above the long-dashed line in the [Mg/Fe] plot (indicated by blue open symbols) are defined as belonging to the high-a population. Stars with data below the long-dashed line (indicated by red filled symbols) are defined as low-a stars. (For interpretation of the references to colour in this figure caption, the reader is referred to the web version of this article.)

There has been recent focus on a series of high-precision abundance studies by Nissen and Schuster (1997) and Nissen and Schuster (2010). They analyzed nearby high-velocity dwarf stars and found that stars with kinematics typical of the halo (VTotal > 180 km s1) and 1.6 < [Fe/H] < 0.4 can be divided into two distinct groups according to elemental abundance trends. One set of stars shows a high level of a elements (‘‘high-a’’ stars) as observed in the canonical halo, while the other set shows a declining trend with increasing iron abundance (‘‘low-a’’ stars). Fig. 17 shows the results for [Mg/Fe] and [a/Fe] versus [Fe/H]. The important aspect of this study is that the analysis contains a large sample of dwarf stars with very similar stellar parameters and adopts a differential approach. The latter ensures that, contrary to most previous studies, the abundance results are internally very robust, with precision for various abundance ratios ranging from 0.02 to 0.04 dex. Nissen and Schuster (2010) also found that the low-a stars tend to have somewhat higher orbital energies than the high-a stars: low-a stars favor retrograde motion and/or extremely large velocities, whereas high-a stars tend to show pro-grade motion. These findings may be consistent with the two-halo picture derived from kinematics and metallicities: low- and high-a stars are expected to belong to the outer and inner halo, respectively. As described below, the low-a stars may represent accreted populations from outside, while the high-a stars may have formed in situ in the halo. Nissen and Schuster (2010) further argued that some of low-a stars may originate from the x Cen progenitor galaxy: the globular cluster x Cen was likely once a nucleus of a dwarf satellite galaxy, where stars from the progenitor galaxy are currently distributed in the Galactic space, including the solar neighborhood (e.g., Dinescu, 2002; Bekki and Freeman, 2003; Mizutani et al., 2003; Meza et al., 2005). If the low-a halo stars indeed originated from this x Cen progenitor galaxy, it would then be conceivable that we should see a substantial number of low-a stars in x Cen. So far, the evidence seems to be the opposite: almost all stars studied to date in x Cen show high a abundances (see e.g., Johnson et al., 2009). In addition, x Cen does not exhibit the characteristic [Na/Fe] split observed by Nissen and Schuster (2010).

It has been speculated that most of these low-a stars belong to the general outer-halo component, which may have been built up from many accreted dwarf galaxies. Thus, it will be intriguing to investigate whether or not the inner and outer halo stars selected in terms of their characteristic kinematics and metallicities actually show a finite difference in [a/Fe] ratios. Ruchti et al. (2010) presented a detailed elemental abundance analysis of 343 red giant and red clump stars selected from RAVE and reobserved at high resolution and high S/N. These stars form a kinematically unbiased, non-local counterpart to the Nissen and Schuster sample (i.e., stars probing the halo and thick disk at some distance from the Sun). Their results show that although several stars with halo kinematics exhibit a low [Mg/Fe] ratio, abundances for the other a elements (e.g., Ca and Si) are not low, in contrast to findings by Nissen and Schuster (2010). Similarly, Ishigaki et al. (2012) conducted an abundance analysis for 97 nearby metal-poor (3.3 < [Fe/H] < 0.5) stars with kinematics characteristic of the thick disk and inner and outer stellar halos based on the high-dispersion spectrograph (HDS) on the Subaru telescope. They showed that for the metallicity range 1.5 < [Fe/H] < 0.5, the thick disk stars show constantly high mean [Mg/Fe] and [Si/Fe] ratios with small scatter and that the inner and outer halo stars show lower mean values for these abundance ratios with larger scatter; [Mg/Fe], [Si/Fe], and [Ca/Fe] ratios for both the inner and outer halo stars also show weak decreasing trends with [Fe/H] >  2. Thus, in both of these studies, no clear confirmation of the results of Nissen and Schuster (2010) was found for the candidate inner and outer halo components. This may be in part because of typical [X/Fe] errors of approximately 0.1 dex (Ruchti et al., 2010; Ishigaki et al., 2012). This error may be still too large to identify the fine abundance details found by Nissen and Schuster (2010). Lindegren and Feltzing (2013) discussed the precision for elemental abundances and how it influences the ability to identify subpopulations in abundance plots as a function of the number of stars studied. Alternatively, there may exist some difference between the kinematically selected sample of Nissen and Schuster (2010) and the non-kinematically selected sample. A more likely reason is that distinction between the inner and outer halo according to their kinematic properties alone is not yet sufficient to incorporate detailed abundance ratios. Specifically, many of retrograde stars with low (U2 + W2)1/2 (<200 km s1) in the Nissen and Schuster (2010) sample may originate from the outer halo, whereas most such stars in the Ishigaki et al. (2012) sample belong to the inner halo according to the definition of the inner/outer halo by Carollo et al. (2010). A better definition of these two halos from detailed kinematics and elemental abundances should be explored. 4.3. Testing the merger picture of halo formation: insight from abundance studies Long after Searle and Zinn (1978) first pointed out the presence of signatures of chaotic merging in the Galaxy formation from elemental abundances of halo globular clusters, evidence of likely relics of past merging of small galaxies were identified. Among these, the discovery of the Sagittarius dSph galaxy (Ibata et al., 1994), which is undergoing tidal interaction with the Galactic potential, is probably one of the most remarkable contributions. This dwarf galaxy also shows a stellar stream, the so-called Sagittarius stream, distributed over the entire sky, as identified by large-area sky surveys (Skrutskie et al., 2006; York et al., 2000). Tidal debris from this stream will be dispersed into the halo space after a few dynamical times and be spatially mixed with field halo stars. SDSS has also identified several other stellar streams or halo substructures, such as the Orphan and Monoceros streams and the Virgo over-density, among others. This is reviewed by Vasily Belokurov in this volume. It has been inferred that elemental abundance distributions for

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these stream stars are distinct from those in field halo stars, although their more detailed elemental abundances are yet to be fully studied (e.g., Chou et al., 2010; Meisner et al., 2012). Many of these halo substructures are found in the outer halo because the dynamical time there is long, of the order of Gyr. This in turn implies that such tidal streams may already have been disrupted and mixed with field stars in the inner halo component. However, while the relaxation time of these substructures is relatively short compared to the life time of the Milky Way, relaxation processes in the phase space defined by integrals of motions such as energy and angular momentum take much longer. Thus, even in the inner halo, evidence of past merging events in the form of kinematic substructures in the phase space can still remain. Helmi et al. (1999) first identified a statistically significant substructure in the angular-momentum space for a nearby metal-poor star sample selected by Beers and Sommer-Larsen (1995) and Chiba and Yoshii (1998), and suggested that this is a remnant of a dSph galaxy that was disrupted during or soon after the Galaxy formed. The presence of this kinematic substructure was later confirmed (Chiba and Beers, 2000; Kepley et al., 2007; Klement et al., 2009; Morrison et al., 2009). Several spectroscopic studies of stars in these kinematic substructures suggest that their abundance patterns appear to be similar to those of general field stars (e.g., Kepley et al., 2007; Ishigaki et al., 2010). Armed with precise astrometric data from Gaia, it will be possible to extract more evidence of past merging events and highlight the merging history of the stellar halo by identifying substructures using the phase-space distribution of stars (Helmi and de Zeeuw, 2000). Measurement of detailed elemental abundance patterns for each such substructure will be essential to obtain a star formation and chemical evolution history for these building blocks of the stellar halo. What about the elemental abundances of stars in the surviving dSph galaxies surrounding and interacting with the Milky Way today? Are they possible remainders of past merging events from the epoch of halo formation? Stars on the red giant branch are bright enough to be studied, even in these distant Galactic satellites. Their [Fe/H] can be measured using the strength of Ca II triplet lines between 840.0 and 870.0 nm (Helmi et al., 2006; Starkenburg et al., 2010). Using a revised calibration for the Ca II triplet, Starkenburg et al. (2010) observed the presence of very metal-poor stars with [Fe/H] < 3 in so-called classical dSph galaxies [bright dSph galaxies known before the recent discovery of much fainter, ultra-faint dwarf (UFD) galaxies]. These results have been confirmed using both Strömgren photometry (Adén et al., 2011) and a detailed analysis of mediumresolution spectra (Kirby et al., 2008). The metallicity distributions for stars in UFDs have been determined; UFDs have stars with [Fe/ H] as low as 3.4 and large spreads in [Fe/H] distributions. Kirby et al. (2008) showed that the mean [Fe/H] ratio in the dSph galaxies continuously decreases with decreasing luminosity, from classical dwarf to UFD galaxies. The latest result for this metallicity–luminosity relation is shown in Fig. 18; the least luminous UFD, Segue I, with MV  1.5, also obeys this relation, with a mean [Fe/H] value of 2.7, although the exact value depends on member assignment (Geha et al., 2009; Norris et al., 2010b). These abundance studies suggest that [Fe/H] distributions for stars assembled from currently surviving Galactic dwarf galaxies are generally in agreement with those for field stars in the Milky Way halo. The implication is that part of the Milky Way halo may have formed from assembly of such dwarf galaxies. However, more detailed abundance patterns for stars in classical dSph galaxies appear to be systematically different from those in the stellar halo (Fig. 19). Shetrone et al. (2001), Shetrone et al. (2003) determined elemental abundances for 36 red giants in seven dSph galaxies and compared these with stars in the Milky

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Fig. 18. Mean metallicity versus absolute magnitude for Milky Way dSph and UFD galaxies (courtesy of Andreas Koch). Data are from Koch (2009), except for the following ultra-faint dwarf galaxies: Adén et al. (2011) for Hercules, Koch et al. (2009) for Bootes II, Grillmair (2009) for Bootes III, Norris et al. (2010b) for Segue 1, Belokurov et al. (2009) for Segue 2, Belokurov et al. (2008) and Walker et al. (2009) for Leo V, and Belokurov et al. (2010) for Pisces II.

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Way, but they found very little to connect the systems. Stars in the dSph galaxies have lower mean [a/Fe] ratios than similar metalpoor Galactic halo stars, whereas the disk stars have much higher metallicities. This finding has been confirmed by spectroscopic studies of classical dSph galaxies (e.g., Venn et al., 2004; Kirby et al., 2009). Venn et al. (2004) studied other elemental ratios including r- and s-process abundances, and found a significant offset in the [Y/Fe] ratio that results in large overabundance of [Ba/Y] in most stars in the dSph galaxies compared to Galactic stars. Thus, the detailed abundance signatures of most of the classical dSph galaxies are distinct from those of the Milky Way halo stars. By contrast, stars in UFD galaxies mainly appear to show similar [X/Fe] ratios to field stars in the Milky Way (e.g., Feltzing et al., 2009; Frebel et al., 2010b; Norris et al., 2010b; Lai et al., 2011) and this is also the case for very metal-poor stars with [Fe/ H] < 3 in the classical dSph galaxies (Cohen and Huang, 2009; Frebel et al., 2010a). In particular, stars in UFD galaxies show a

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dependence of [C/Fe] on [Fe/H] similar to that seen in Galactic field halo stars (Norris et al., 2010b; Frebel et al., 2010b; Gilmore et al., 2013). It is also worth noting that the relative abundances of a large number of elements in UFDs are very similar to those in the stellar halo (Frebel et al., 2010b; Norris et al., 2010a). Thus, it has been suggested that such UFD galaxies were indeed building blocks for very metal-poor stars in the stellar halo. 4.4. Chemodynamical evolution of the stellar halo The paucity of metals in halo stars compared to disk stars indicates that the stellar halo component formed when interstellar gas was still metal-poor prior to disk formation. Several chemical evolution models have been put forward to explain the metallicity distribution observed, which has a peak at [Fe/H]  1.6 and a tail at the metal-poor end. These models include a simple model of chemical evolution including gaseous outflow (Hartwick, 1976) or outflow plus an additional early inflow (Prantzos, 2003). While these models assume that stars form from well-mixed gas clouds, a proper account of incomplete mixing of the interstellar medium is required to explain the abundance properties of extremely metalpoor stars. Indeed, inhomogeneous chemical evolution models developed by Tsujimoto and Shigeyama (1998), Ishimaru and Wanajo (1999), and Argast et al. (2000), for example, successfully reproduce the abundance scatter observed for halo stars with metallicities below 2.0 dex. As described above, halo stars are also characterized by large random motions with an anisotropic velocity distribution, depending on their galactocentric distance and metallicity (Kafle et al., 2012). This suggests that combined analysis taking into account both dynamical and chemical evolution processes in association with the formation of the Milky Way is required to fully understand the basic properties of the stellar halo. Since Eggen et al. (1962), many numerical studies of galaxy formation have been performed. In particular, Larson (1969, 1974) conducted a series of numerical simulations of the collapse of an initially gaseous proto-galaxy and concurrent transformation of gas into stars, including chemical evolution and supernova explosion, in which the monolithic dissipative collapse hypothesis proposed by Eggen et al. (1962) was adopted. Remarkable progress was made by Katz and Gunn (1991), who first simulated dissipational galaxy formation in the context of hierarchical clustering theory, using the three-dimensional TREESPH method. They demonstrated that rotation-dominated disk galaxies can form from initial density perturbations embedded in the Hubble flow, but they did not explore the origin of the stellar halo component. Later, using state-of-art chemodynamical codes in the context of CDM models, Steinmetz and Mueller (1994, 1995) showed the spatial distribution of stellar populations of different ages, including disk younger stars in the disk surrounded by diffusely distributed old, metal-poor halo stars. Following the Hipparcos catalogue (ESA, 1997) Bekki and Chiba (2000), Bekki and Chiba (2001) conducted the first detailed comparison of a chemodynamical simulation of galaxy formation based on CDM models with observed halo kinematics and abundances. They showed that galaxy formation in CDM models includes aspects of both the Eggen et al. (1962) and Searle and Zinn (1978) scenarios, dissipational collapse and dissipationless merging in a proto-galactic system, which eventually provides the kinematics and elemental abundances observed for the halo stars. In particular, they succeeded in reproducing the orbital eccentricity distribution observed for halo stars near the Sun, which is key to understanding the early evolution of the Milky Way halo (Hattori and Yoshii, 2011). More detailed calculations for the formation of a stellar halo in a Milky-Way-sized dark matter halo have been made based on hierarchical assembly of CDM halos. Bullock and Johnston (2005)

Fig. 20. Stellar density and metallicity profiles simulated by Font et al. (2011). Dotted and solid black curves represent all stars and stars in the spheroid only, respectively. Dashed blue and dot-dashed red curves represent accreted halo stars and spheroid stars that formed in situ, respectively. Note that in situ star formation contributes significantly to the stellar halo at r < 20 kpc and that stars formed in situ at r < 50 kpc have higher metallicity than those originating from accreted satellites. (For interpretation of the references to colour in this figure caption, the reader is referred to the web version of this article.)

adopted a hybrid semi-analytic plus N-body approach to distinguish between the evolution of light and dark matter in accreted satellites. They generated 11 random realizations of hierarchically merged, Milky-Way-sized dark matter halos. Their model accounts for star formation processes in each sub-halo. Their stellar halos formed entirely from accreted satellites have density profiles that are consistent with observations, in the form of power laws at r > 30 kpc with q / ra and a ’ 3–4. They also showed that stellar halos contain numerous substructures as relics of past accretion events. This is indeed seen in many nearby galaxies (Martínez-Delgado et al., 2010). More recently, Cooper et al. (2010) used a very high-resolution N-body simulation of CDM models and showed how more detailed structures of stellar halos form. What mechanism in the galaxy formation process is responsible for the two-component structure of the stellar halo seen in the Milky Way? As envisaged by, for example, Bekki and Chiba (2001), for example, a complete picture of the stellar halo formation involves not only dissipationless accretion/merging of satellites but also dissipative processes leading to collapse and star formation. Indeed, recent numerical simulations provide more detailed and realistic insight into this in situ formation process of halo stars and the contribution of accreted stars from outside in the dark matter halo of the host (Zolotov et al., 2009, 2010; Oser

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et al., 2010; Font et al., 2011; McCarthy et al., 2012; Tissera et al., 2012, 2013). Zolotov et al. (2009, 2010) investigated the competing importance of in situ star formation and accretion of stars from sub-halos. They showed that all four simulated galaxies are surrounded by a stellar halo whose inner regions (r < 20 kpc) contain both accreted stars and a stellar population formed in situ. The outer regions of the halos were assembled through pure accretion and disruption of satellites. This two-step process of halo formation was further studied based on a state-of-art simulation code by Font and her collaborators (Font et al., 2011; McCarthy et al., 2012). They found that a flattened inner halo component possibly originates from the in situ star formation mode, whereas accreted halo stars dominate the outer part of the stellar halo, as shown in the upper panel of Fig. 20 (Font et al., 2011). The lower panel of Fig. 20 demonstrates that stars formed in situ have a higher metallicity than those originating from accreted satellites, as suggested by recent observations.

5. Outlook We have presented a snapshot of current knowledge of the chemical make-up of the Milky Way and, by combining this information with kinematic information, gave a picture of the chemodynamical status of the disk, bulge, and halo of the Milky Way. However, this knowledge barely scratches the surface of the immensely rich stellar systems that make up the Milky Way. For example, studies that found differences between the stellar disks are based on small samples of a few hundred to a thousand stars in the direct solar neighborhood (within 100 pc). We have no detailed elemental abundances data for stars between us and the Galactic bulge. Thus, a full 8000 pc represents a virtual knowledge desert for which we only have a minimum of information about the properties of the stellar disk(s). Elucidation of the formation and evolution of galaxies in a cosmological context requires an understanding of the stellar disks, so this lack of knowledge is a severe handicap to any further progress. Nonetheless, observational results reviewed here based on the assembly of basic stellar spectra will be largely extended by several ongoing spectroscopic projects, such as Gaia-ESO Survey (Gilmore et al., 2012), APOGEE, HERMES, and RAVE. Furthermore, in synergy with the Gaia mission, the planned next-generation spectroscopic surveys using massively multiplexed fiber-fed spectrographs, including LAMOST (Zhao et al., 2012), WHT/WEAVE (Balcells et al., 2010), Subaru/PFS (Sugai et al., 2012), VISTA/4MOST (de Jong et al., 2012), VLT/MOONS (Cirasuolo et al., 2012), and ngCFHT (Côté, 2012), will have a significant impact on Milky Way studies. The combination of distances and proper motions from Gaia and these ground-based, massive follow-up surveys will allow us to analyze the detailed six-dimensional phase–space distribution of Galactic stars and their distribution in multi-dimensional abundance and age space. Taking such insights together with further and deeper understanding of basic galaxy formation processes based on theoretical modeling, we should be able to construct an ultimate picture of Galaxy formation in the near future.

Acknowledgments Work by S.F. is partly supported by Grant No. 2011-5042 from the Swedish Research Council. M.C. was supported in part by a Grant-in-Aid for Scientific Research (20340039, 18072001) from the Ministry of Education, Culture, Sports, Science and Technology of Japan and by the JSPS Core-to-Core Program ‘‘International Research Network for Dark Energy’’.

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