The Astrophysical Journal, 771:67 (25pp), 2013 July 1

doi:10.1088/0004-637X/771/1/67


C 2013. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

CHEMICAL ABUNDANCES OF THE MILKY WAY THICK DISK AND STELLAR HALO. II.
SODIUM, IRON-PEAK, AND NEUTRON-CAPTURE ELEMENTS
M. N. Ishigaki1 , W. Aoki2 , and M. Chiba3

1 Kavli Institute for the Physics and Mathematics of the Universe (WPI), University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8583, Japan;

miho.ishigaki@ipmu.jp
2 National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan; aoki.wako@nao.ac.jp
3 Astronomical Institute, Tohoku University, Aoba-ku, Sendai 980-8578, Japan; chiba@astr.tohoku.ac.jp

Received 2012 November 21; accepted 2013 May 2; published 2013 June 17

ABSTRACT
We present chemical abundance analyses of sodium, iron-peak, and neutron-capture elements for 97 kinematicallyselected thick disk, inner halo, and outer halo stars with metallicities −3.3 < [Fe/H] < −0.5. The main aim of thisstudy is to examine chemical similarities and differences among metal-poor stars belonging to these old Galactic
components as a clue to determine their early chemodynamical evolution. In our previous paper, we obtained
abundances of α elements by performing a one-dimensional LTE abundance analysis based on the high-resolution(R ∼ 50,000) spectra obtained with the Subaru/HDS. In this paper, a similar analysis is performed to determine
abundances of an additional 17 elements. We show that, in metallicities below [Fe/H] ∼ −2, the abundance
ratios of many elements in the thick disk, inner halo, and outer halo subsamples are largely similar. In contrast, in
higher metallicities ([Fe/H]  −1.5), differences in some of the abundance ratios among the three subsamples are
identified. Specifically, the [Na/Fe], [Ni/Fe], [Cu/Fe], and [Zn/Fe] ratios in the inner and outer halo subsamples
are found to be lower than those in the thick disk subsample. A modest abundance difference between the two halosubsamples in this metallicity range is also seen for the [Na/Fe] and [Zn/Fe] ratios. In contrast to that observedfor [Mg/Fe] in our previous paper, [Eu/Fe] ratios are more enhanced in the two halo subsamples rather than in
the thick disk subsample. The observed distinct chemical abundances of some elements between the thick disk andinner/outer halo subsamples with [Fe/H] > −1.5 support the hypothesis that these components formed throughdifferent mechanisms. In particular, our results favor the scenario that the inner and outer halo components formed
through an assembly of multiple progenitor systems that experienced various degrees of chemical enrichments,
while the thick disk formed through rapid star formation with an efficient mixing of chemical elements. The lower[Na/Fe] and [Zn/Fe] observed in stars with the outer halo kinematics may further suggest that progenitors withlonger star formation timescales contributed to the buildup of the relatively metal-rich part of stellar halos.
Key words: Galaxy: formation – Galaxy: halo – stars: abundances
Online-only material: color figures, machine-readable tables
disk stars are known to be characterized by higher [α/Fe]
ratios than the thin disk stars (Fuhrmann 1998; Bensby et al.
2003, 2005; Reddy et al. 2003, 2006; Prochaska et al. 2000;
Ruchti et al. 2011), more similar to the stars in the Galactic
bulge (Melendez et al. 2008; Bensby et al. 2010). Systematic
analyses of low-resolution spectra for a large number of stars
from the Sloan Extension for Galactic Understanding andExploration (SEGUE) have significantly advanced our viewof this component (e.g., Carollo et al. 2010; Lee et al. 2011;
Cheng et al. 2012). In particular, based on the SDSS/SEGUE
data, Carollo et al. (2010) suggest that the previously known
metal-poor tail (−2.7 < [Fe/H] < −1.0) of the metallicity
distribution function of the thick disk, which is called the“metal-weak thick disk” (MWTD; e.g., Beers et al. 2002), could
indeed be associated with an independent stellar component. On
the other hand, Bovy et al. (2012) questioned the distinct natureof the thick disk apart from the thin disk based on the analysisof chemically defined subpopulations by carefully taking into
account possible selection biases in the SDSS spectroscopicdata.
The MW stellar halo is known as another oldest remnantof the early chemodynamical evolution of our Galaxy. Recent
large surveys such as SDSS reveal that the MW stellar halo
is highly structured and cannot be approximated by a smooth
spherical distribution of metal-poor stars (e.g., Ivezić et al.
2012). It became clear that the stellar halo contains a numberof substructures in spatial distribution (e.g., Ibata et al. 1994;

1. INTRODUCTION
Due to the advent of dedicated photometric and spectroscopic
surveys such as the Sloan Digital Sky Survey (SDSS), theold metal-poor components of our Milky Way (MW) Galaxy,
namely, the thick disk and the stellar halo, are found to be
far more complex than previously thought exhibiting various
degrees of substructures. The origins of the complex natureof these old structural components have been studied throughobservations and theoretical models as one of the central issues
to unveil how the MW formed and evolved along with theevolution of galaxies in the universe. Although these efforts
were successful in greatly advancing our view of these old MW
components, their origins remain unclear.
Although the MW thick disk has been known for many years,
its dynamical and chemical structures are still controversial inconstraining its formation mechanism. The MW thick disk was
originally discovered as an extra component required to fit thestellar number density distribution over the thin disk component
at distances from the disk plane greater than ∼1 kpc (Yoshii
1982; Gilmore & Reid 1983). The thick disk is widely believed
to have been formed at an early epoch of Galaxy formation
because constituent stars have older ages and lower metallicities
compared to thin disk stars (e.g., Bensby et al. 2003). In thesolar neighborhood, thick disk stars lag in rotational velocities
behind thin disk stars by ∼20–50 km s−1 (Chiba & Beers
2000; Carollo et al. 2010). Chemical abundances of the thick1

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba

Fulbright 2002; Roederer 2008; Zhang et al. 2009; Ishigaki
et al. 2010; Nissen & Schuster 2010, 2011; Ishigaki et al. 2012,
hereafter Paper I). Stephens & Boesgaard (2002) reported thatstars having a large apocentric distance (Rapo ) tend to havelower [α/Fe] while no correlations of abundances on other
orbital parameters were identified. Since sample stars withextreme outer halo kinematics were still limited, the presenceor absence of the correlation remained uncertain. Recently,
Nissen & Schuster (2010) suggest that their sample of nearby
dwarf stars can be divisible into two distinct groups in terms ofabundances of several α elements, namely, the low-α and high-αstars. The two chemically distinct groups of stars tend to show
different characteristic kinematics. Nissen & Schuster (2011)
further reported that these groups show different abundancesin other elements such as Na and Zn. These studies imply
that the stars more metal-poor than [Fe/H] < −0.5 cannot be
formed within a single well-mixed gas but more likely formed
in different pre-Galactic clumps that have their own chemicalenrichment history.
It remained uncertain, however, how the presence of thechemically distinct groups of stars in the solar neighborhood fits
into the formation scenario for thick disk and inner/outer stellar
halos. A sample of stars with a wide range of kinematics andmetallicities, including those characteristics of the thick disk,
inner halo, and outer halo components, would provide useful
insights into this issue. In this paper, we investigate similarities
and differences in detailed chemical abundance patterns among
the kinematically selected thick disk, inner halo, and outer halostars. This allows us to investigate whether or not the stars
in these components with their overlapping metallicity range
formed under the influence of a similar chemical enrichment
history. For this purpose, a sample of 97 metal-poor ([Fe/H] <−0.5) stars spanning a wide range of orbital parameters and[Fe/H] are studied. In Paper I, we present the abundance analysis
of α-elements (Mg, Si, Ca, and Ti) for the sample stars. We
showed that in a metallicity range of −1.5 < [Fe/H] < −0.5,
kinematically defined thick disk stars have higher [Mg/Fe] and[Si/Fe] with small scatter, while the inner and outer halo stars
show lower average abundance ratios for these elements withlarger scatter. In the present study, we further present the results
for sodium, iron(Fe)-peak, and neutron-capture elements forthe same sample of stars and investigate possible scenarios thatconsistently explain the observed abundance patterns in each of
the old Galactic components.
This paper is organized as follows. In Section 2, we present asummary of our sample stars and their membership to the thick
disk, inner halo, and outer halo subsamples that have been takenfrom Paper I. In Section 3, a brief review of the observation,
which was fully described in Paper I, is presented. Then,
determination of stellar atmospheric parameters and abundance
analyses of individual elements are described. In Section 4,
we present the abundance results of the distribution in the[X/Fe]–[Fe/H] plane for each of the kinematically selectedsubsamples and examine the correlation between [X/Fe] andorbital parameters. Finally, Section 5 discusses the interpretationof the abundance differences and similarities among the threesubsamples and their implications for the formation of thesecomponents.

Newberg et al. 2002; Majewski et al. 2003; Jurić et al. 2008)
or kinematics (e.g., Helmi et al. 1999; Chiba & Beers 2000;
Kepley et al. 2007; Schlaufman et al. 2009) of constituent
stars. The most prominent example is a currently merging
dwarf galaxy, Sagittarius, and its tidal debris (e.g., Helmi 2008;
Ivezić et al. 2012 and references therein). Furthermore, Carollo
et al. (2007, 2010) recently confirmed that the MW stellar halo
is divisible into two globally overlapping stellar components,
namely, the inner and outer stellar halos, based on the analysisof kinematics and metallicity for a large sample of calibrationstars obtained by SDSS/SEGUE. Those authors showed that the
inner halo dominates at Galactocentric distances r smaller than
∼10–15 kpc. It has a modestly flattened distribution of stars
with a mean metallicity of [Fe/H] ∼ −1.6. The kinematics ofthe stars likely belonging to the inner halo are characterized bya zero to slightly prograde mean rotational velocity with large
velocity dispersions. In contrast, the outer halo was found to
dominate at r > 10–15 kpc and exhibit a more spherical stellar
density distribution. The mean metallicity was estimated to be[Fe/H] ∼ −2.2, which is lower than that of the inner halo.
The kinematics of the outer halo component are characterized
by a larger number of extreme motions (e.g., large progradeor retrograde orbit) that are the outlier of the typical inner halostars (Carollo et al. 2007, 2010). Although the definite rotational
properties of the two components remain under discussion (e.g.,Schönrich et al. 2011; Beers et al. 2012), the proposed differencein the structural properties between the inner and outer partsof the MW halo suggests that these two components formed
through different mechanisms (Carollo et al. 2010).
Chemical abundances of individual metal-poor stars that
constitute the thick disk and stellar halos provide a unique
opportunity to test theoretical models for the chemodynamical
evolution of these components, particularly for the possible
progenitors of these systems (Freeman & Bland-Hawthorn
2002). It has long been known that chemical abundances ofmetal-poor stars in the solar neighborhood are characterized byan enhancement of α elements (Mg, Si, Ca, and Ti) to ironabundance ratios ([α/Fe]; e.g., Luck & Bond 1981; McWilliam
et al. 1995; Cayrel et al. 2004; Lai et al. 2008). This result is
usually interpreted as evidence that the MW field metal-poorstars are formed out of gas enriched mainly through Type IIsupernovae (SNe) of massive stars, which would yield high[α/Fe] (e.g., Matteucci & Greggio 1986). On the other hand,observations of spatial distribution, kinematics, and metallicity
for a number of MW stars suggest that the metal-poor stars in the
Galaxy are a mixture of apparently different stellar populations(e.g., Bell et al. 2008). If each of these populations has their own
chemical enrichment histories, we would expect to distinguishthese populations through kinematics and detailed chemical
abundances of individual metal-poor stars.
In this context, any correlation between orbital kinematics andchemical compositions is of particular interest since a star withextreme kinematics is more likely to have been accreted froman external system, which may have quite different chemicalenrichment histories in their birthplace. Detailed chemical
abundances of such stars provide us with clues for understanding
properties such as star formation rates, efficiency of galactic
wind, initial mass function (IMF), etc., of possible building
blocks of our Galaxy (Lanfranchi & Matteucci 2003). Taking
advantage of a high-resolution spectrograph on 8–10 m classtelescopes, comprehensive studies on the correlation between
detailed chemical abundances and kinematics of nearby halostars have been carried out (Stephens & Boesgaard 2002;

2. THE SAMPLE
The sample of 97 dwarf and giant stars with [Fe/H]  −0.5
was selected from the catalogs of Carney et al. (1994), Ryan &Norris (1991), and Beers et al. (2000) based on their kinematics.
2

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba

The proper motions and distance estimates were partly updated
from those in the original catalog as described in Paper I.
The radial velocities for the sample stars were also updated
to those measured from the high-resolution spectra obtainedin our observation. The orbital parameters Rapo (apocentric
distance), Zmax (maximum distance from the Galactic plane),and e (orbital eccentricity) were calculated by adopting theStöchel-type Galactic potential in the same manner as describedin Chiba & Beers (2000). More details on the sample selection
and kinematics of the sample stars are described in Paper I.

component was reported to dominate in the metallicity range−1.0  [Fe/H]  −0.4 (e.g., Wyse & Gilmore 1995), while
our thick disk subsample extends to the metallicity as low as−2.7. At metallicities below [Fe/H] < −1.0, Carollo et al.
(2010) suggested that the independent MWTD component isrequired to account for the observed [Fe/H] and Vφ distribution
for their sample stars close to the Galactic plane. They also
reported that the metallicity of the MWTD component is in arange −1.7 < [Fe/H] < −0.8 and Vφ ∼ 100–150 km s−1 with adispersion of ∼35–45 km s−1 . Four of the thick disk stars in our
sample with [Fe/H] < −0.8 have rotational velocities similar
to the MWTD (140 < Vφ < 190 km s−1 ), where two of themost metal-poor stars have the lowest Vφ values. It is unclear
whether these two stars represent the lowest metallicity tail ofthe MWTD or interlopers from the halo component.
Chemical abundances other than iron for the MWTD havebeen investigated by several studies, in which a distinct chemicalsignature for this component was not clearly identified. Reddy& Lambert (2008) studied elemental abundances for α, ironpeak, and neutron-capture elements in the 14 candidate MWTD
stars and found that their abundances are indistinguishable fromhalo stars with similar metallicity. Ruchti et al. (2011) studied
abundances of iron and α-elements for a large sample of metalpoor stars based on the medium-resolution spectroscopic datafrom the Radial Velocity Experiment. They reported that themetal-poor thick disk stars are enhanced in [α/Fe] ratios similar
to the halo stars.
In Paper I, we investigate [α/Fe] for the four stars with similar
properties as the elusive MWTD as mentioned above. As aresult, these stars were shown to have higher [Mg/Fe] or [Si/Fe]
ratios than the inner/outer halo stars. Although the number ofstars in this sample is very small to extract a definite conclusion
about the properties of the proposed MWTD, we later compare
their chemical abundances other than α elements with those of
the typical thick disk, inner halo, and outer halo stars in our
sample.
Our classification of inner and outer halo stars may not be
representative of the stellar halo components observed in othersurveys (Carollo et al. 2010; De Jong et al. 2010; An et al.
2013). Stars with outer-halo-like kinematics are classified asouter halo members regardless of their [Fe/H]. Following thesame criteria, stars with inner halo kinematics are assigned tothe inner halo component independently from their metallicity.
As a result, our sample of inner halo and outer halo stars bothspan the wide range in metallicity −3.5 < [Fe/H] < −0.4. This
is in contrast to the inner/outer halo division reported by Carollo
et al. (2010), where the inner halo stars have a peak metallicityof −1.6 while the outer halo stars have ∼−2.2 dex. Note thatsome of the stars assigned to the inner halo or the outer halo
have metallicity in agreement with these components as derived
in previous works.
The reason for this metallicity difference is not clear. One possible explanation is that our sample selection in the solar neighborhood (<1–2 kpc) may be biased toward/against particular
metallicity among the inner halo or outer halo components. In
the following, we focus on comparing abundance ratios between
kinematically defined subsamples at a given metallicity.

2.1. Kinematics and Membership Assignment
Based on the kinematics, we assigned the membership for the
thick disk, inner halo, or outer halo components to each of thesample stars as described in Paper I. In short, we calculated
the probabilities that each of the sample stars belongs to thethick disk (PTD ), inner halo (PIH ), or outer halo components(POH ), based on their space velocities in the Galactic cylindricalcoordinate (VR , Vφ , and VZ ). In this calculation, the mean
velocities and dispersions for the thick disk, inner halo, andouter halo components as well as a fractional contribution ofeach component at different Zmax were adopted from the values
obtained by Carollo et al. (2010) based on their analyses of
SDSS DR7. Then, the thick disk, inner halo, and outer halostars were defined as the stars with PTD > 0.9, PIH > 0.9, and
POH > 0.9, respectively. Other stars with PTD , PIH , POH  0.9
were classified as either the thick disk/inner halo or the inner
halo/outer halo intermediate populations. In the remainder ofthis paper, we conventionally refer to the stars classified as the
thick disk, inner halo, and outer halo categories as the “thick
disk, inner halo, and outer halo stars/subsamples,” respectively.
In the above definitions, 12, 34, and 37 stars are assigned to thethick disk, inner halo, and outer halo components, respectively.
The adopted criteria for the membership assignment are purely
based on kinematics while metallicities for each component are
not taken into account. We briefly describe its consequence inthe next subsection.
As shown in Figure 1 of Paper I, the thick disk stars in our
sample have a mean rotational velocity of Vφ ∼ 180 km s−1 ,whose orbit is confined to ∼1 kpc above and below the Galactic
plane (Zmax < 1 kpc). The inner halo stars show no rotationVφ ∼ 0 on average and exhibit larger velocity dispersion.
Finally, the outer halo stars show a much larger dispersion in Vφand some have extreme prograde or retrograde rotation. At thesame time, almost all stars with orbits that reach the distance of>10 kpc from the Galactic plane were classified as outer halostars by definition.
2.2. Metallicity of the Sample
Our classification of the sample stars into the thick disk,
inner halo, and outer halo subsamples described above is purely
based on kinematics. As a result, metallicities for each of thethree subsamples may be different from those of the thick disk,
inner halo, and outer halo components obtained in previous
works (e.g., Carollo et al. 2010). In this subsection, we compare
metallicities of our thick disk, inner halo, and outer halosubsamples with the metallicity distribution functions reported
in the literature for each Galactic component.
The thick disk stars in our sample span a lower metallicity
range than that reported for the canonical thick disk component
and likely include stars with chemical and kinematical properties
similar to the MWTD component. The canonical thick disk

3. OBSERVATION AND DATA ANALYSIS
3.1. Observation and Data Reduction
The observations for all of the sample stars were madewith the High Dispersion Spectrograph (HDS; Noguchi et al.
3The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba
Table 1
Equivalent Widths

Object Name

Z/Ion

Elementλ(Å)

log gf
(dex)

χ
(eV)

EW
(mÅ)

Flaga

Refs.bhfs

BD+01◦ 3070
BD+01◦ 3070
BD+01◦ 3070
BD+01◦ 3070
BD+01◦ 3070

11 1
21 2
21 2
21 2
21 2

Na i
Sc ii
Sc ii
Sc ii
Sc ii

5682.634400.40
4670.425031.025239.82

−0.70
−0.54
−0.72
−0.40
−0.77

2.10
0.61
1.36
1.36
1.45

9.82
67.22
35.52
45.74
26.14

1
1
1
1
1

NS10
I06
R10
I06
I06

hfs
hfs
hfs
hfs

Notes.
a 1: used in the abundance analysis; 0: not used in the abundance analysis.
b Reference of the adopted log gf . A complete list of references is given in the electronic version of this table.
(This table is available in its entirety in a machine-readable form in the online journal. A portion is shown here
for guidance regarding its form and content.)

2002) mounted on the Subaru telescope during 2003–2010. The
wavelength range of ∼4000–6800 Å was covered with a spectral
resolution of R ∼ 50,000 for most of the sample stars, while
the several sample stars (G 64−12, G 64−37, BD+13 2995,
G 14−39, and G 20−15) were observed with R ∼ 90,000. The
data reduction including bias correction, cosmic-ray removal,
flat fielding, scattered light subtraction, wavelength calibration,and continuum normalization was performed with standard
IRAF routines. Details of the observational setting and their
reduction procedures are described in Ishigaki et al. (2010) andPaper I. The equivalent widths (EWs) of absorption lines were
measured by fitting Gaussian to each feature.

HD 171496 from Alves-Brito et al. (2010), in which the same
Teff -scale of Ramı́rez & Meléndez (2005) was used but with amore sophisticated E(B − V ) estimate based on the Na i D line.
For LP 751−19, we adopt the value from Paper I. As described
below, the LP 751−19 shows anomalous abundances for someelements compared to the behavior of other sample stars, whichmight result from adopting a wrong Teff value. The log g and ξ
values of HD 171496 are updated, adopting the revised value ofTeff , based on the Fe i/Fe ii excitation equilibrium and the Fe i
abundance–EW relation.
3.2.2. Abundances

We use the EWs of metal absorption lines in the derivationof abundances for most of the elements. The measured EWs
are given in Table 1. For Cu and Eu, we adopt a spectral
synthesis for their abundance estimates. The derived abundances
are normalized with the solar values from Asplund et al. (2009)
to obtain the [X/H] value. The [X/Fe] ratios are then derived
by normalizing [X/H] with [Fe i/H] or [Fe ii/H] for neutral
or ionized species, respectively. The derived abundances andadopted stellar atmospheric parameters are given in Table 2.
We describe below notes on derivation of individual elemental
abundances and atomic data.
Sodium. Sodium abundance is mainly determined from the
Na i lines at 5682.6, 6154.2, and 6160.8 Å. We avoid using
the Na i resonance lines at 5890/5896 Å since large negative
non-LTE correction (log NLTE − log LTE ) up to ∼−0.5 dex
was previously reported for these lines (Takeda et al. 2003;
Andrievsky et al. 2007). On the other hand, for the Na i lines usedin the present analysis, the non-LTE calculation by Takeda et al.
(2003) suggests that its correction is not more than −0.11 dex
for their sample of modestly metal-poor (−1.0 < [Fe/H] < 0.0)
dwarf stars. Since the reported amount of correction is not
significantly larger than the errors in the Na abundances inthis study, we shall consider that correction for the dwarf stars
in our sample within this metallicity range is negligible. The
values for the correction may vary depending on Teff , log g, or
metallicities in a complex way. However, we assume that thecorrection is small for our whole sample of stars and simply
adopt the abundances derived from the LTE analysis with nocorrection.
Scandium. Sc abundances have been determined from EWsof Sc ii lines. The hyperfine splitting (hfs) was taken into
account in the abundance derivation, adopting the wavelengthand the fractional strength of each hyperfine component fromthe Kurucz (1995) database. The total log gf values for each line
are normalized to those in Ivans et al. (2006) and Roederer et al.

3.2. Abundance Analysis
Abundance analyses are performed by using an LTE codewith model atmospheres of Castelli & Kurucz (2003), which
is widely described in Aoki et al. (2009) and in Paper I. In
this subsection, we describe additional details on derivation ofindividual elemental abundances.
3.2.1. Stellar Atmospheric Parameters

We basically adopt the effective temperature (Teff ), surface
gravity (log g), and micro-turbulent velocity (ξ ) that were
estimated and used in Paper I. The Teff was estimated by thecolor Teff relation using the calibrations of Casagrande et al.
(2010) for dwarfs and Ramı́rez & Meléndez (2005) for giants
that are based on the infrared flux method. As mentioned inPaper I, adopting Teff from the color results in a non-negligible
slope in the iron abundances versus excitation potentials of
the Fe i lines for some of the sample stars. The two stars
HD 171496 and LP 751−19 show exceptionally large slopes(0.09 and 0.14 dex eV−1 , respectively) compared to the medianvalue of −0.06 dex eV−1 for the whole sample. A possible
reason for the peculiar behavior of these stars is that theE(B − V ) values may be underestimated. Both HD 171496 andLP 751−19 are located at the Galactic latitude, b = −7.7317
and −5.2902, respectively. For these directions, the E(B − V )values in Schlafly & Finkbeiner (2011) have been reported tobe 0.38 and 0.61, respectively. On the other hand, we estimateE(B − V ) = 0.10 and 0.025, respectively, using the iterativealgorithm to take into account the finite distance to each star (294
and 56 pc, respectively). These estimates may be affected by theuncertainty in the distance estimate as well as the uncertainty
in the distribution of dust near the Galactic disk. Since accurate
estimates of E(B − V ) from interstellar Na i lines is difficult
for the spectral resolution of our data, we adopt the Teff value of
4

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba
Table 2
Atmospheric Parameters and Abundances

Object Name

Teff
(K)

log g
(dex)

ξ
(km s−1 )[Fe i/H](dex)[Fe ii/H](dex)[Na/Fe](dex)[Sc/Fe](dex)[V/Fe](dex)[Cr i/Fe](dex)[Cr ii/Fe](dex)

BD+01◦ 3070
BD+04◦ 2466
BD+04◦ 2621
BD+09◦ 2870
BD+10◦ 2495

5404
5223
4754
4632
4974

3.65
2.02
1.63
1.30
2.29

1.18
1.72
1.72
1.63
1.64

−1.36 ± 0.14
−1.93 ± 0.14
−2.40 ± 0.16
−2.38 ± 0.17
−2.01 ± 0.15

−1.37 ± 0.13
−1.94 ± 0.12
−2.41 ± 0.12
−2.42 ± 0.12
−2.01 ± 0.12

−0.30 ± 0.12
−0.08 ± 0.14
−9.99 ± 0.00
−9.99 ± 0.00
−9.99 ± 0.00

0.18 ± 0.07
0.13 ± 0.08
0.08 ± 0.08
0.07 ± 0.08
0.02 ± 0.07

0.34 ± 0.10
−9.99 ± 0.00
−0.03 ± 0.12
−0.14 ± 0.11
−0.04 ± 0.10

−0.12 ± 0.06
−0.20 ± 0.08
−0.24 ± 0.09
−0.31 ± 0.08
−0.21 ± 0.07

0.14 ± 0.08
0.28 ± 0.10
0.17 ± 0.07
0.11 ± 0.07
0.17 ± 0.07

(This table is available in its entirety in a machine-readable form in the online journal. A portion is shown here for guidance regarding its form and content.)(2010). The effect of the hfs is very small, which is typically∼0.03 dex or less, for the present sample.
Vanadium, chromium, nickel, zinc, yttrium, and zirconium.
The EWs of the V i, Cr i, Cr ii, Ni i, Zn i, Y ii, and Zr ii lines wereused for their abundance determination, where log gf values
were mainly adopted from Ivans et al. (2006) and Roederer
et al. (2010). For the Cr i lines, the log gf values from therecent laboratory measurements of Sobeck et al. (2007) are also
included.
Manganese. For the abundance determination of Mn, EWs of
Mn i lines were used. The log gf values of Ivans et al. (2006) andRoederer et al. (2010) and the new measurements of BlackwellWhitehead & Bergemann (2007) were adopted. hfs of theselines is taken into account based on the fractional strengths ofeach component in the Kurucz (1995) database.
Cobalt. Co i lines are used for the abundance analysis of Co
taking into account the hyperfine structure for these lines. The
log gf values and the atomic data for the hyperfine structure
were taken from Pickering (1996).
Copper. Abundances of Cu have been obtained from the
Cu i line at 5105.5 Å for a subset of the sample stars. We
have employed a spectral synthesis for the abundance estimate,since line broadening due to hyperfine and isotopic splitting is
expected for this line (Simmerer et al. 2003). The line splitting
and the fractional strength of each hyperfine component wereadopted from the Kurucz (1995) database, while the overalllog gf value for this line was taken from Fuhr & Wiese (2005).
The isotopic fractions of the two stable isotopes, 63 Cu and 65 Cu,
were assumed to be the solar-system fractions of 69% and 31%,respectively (Simmerer et al. 2003).
Strontium. For most of the sample stars, the strong Sr ii lines
at 4077.2 Å and 4215.5 Å are identified. However, these lines
are saturated or blended especially in the sample stars with
[Fe/H]  −1.0 and are not useful for the abundance estimates(a change in the Sr abundance makes only a small change inthe line strength). We exclude the sample stars with equivalent
widths log(EW/λ)  −4.6 in the following discussions on the
Sr abundances. The log gf values for these lines were adopted
from Fuhr & Wiese (2005).
Barium. Barium abundances have been obtained from EWsof the Ba ii lines at 4554.0, 4934.1, 5853.7, and 6141.7 Å. In
the abundance calculation, hfs and isotopic shifts for each of thefive stable Ba isotopes (134 Ba, 135 Ba, 136 Ba, 137 Ba, and 138 Ba)
have been taken into account. The fractional contribution of eachhyperfine component and their wavelength shifts are basicallyadopted from McWilliam (1998), with some modifications
described below.
The assumption about the isotopic fraction, which is determined by the fractional contribution of the s- and the r-process insynthesizing the observed Ba, affects the abundance determination, since the odd mass number isotopes show hfs while the evenmass number isotopes do not. In the present analysis, we adopt
the isotopic fraction expected for the solar-system r-processcomponent from McWilliam (1998), except for the two Ba-richstars. As shown later, these two stars show exceptionally low[Eu/Ba] ratios, which indicate a significant contribution of thes-process nucleosynthesis. For these two stars, we assume thesolar s- and r-process mix (81% and 19%, respectively) predicted by Arlandini et al. (1999) and the isotopic fractions of
the s-process component from Anders & Grevesse (1989). Assuming the solar-system r-process isotopic fraction may not beadequate for the sample stars with [Fe/H] > −1.5, since some
contribution from the s-process is expected in this metallicity
range. In order to reduce the abundance errors due to the uncertainty in the isotopic fraction, we exclude the two resonancelines at 4554.03 Å and 4934.10 Å, which are particularly sensitive to the assumed isotopic fractions in the Ba abundance
determination for the sample stars with [Fe/H] > −1.5. Other
two lines in the redder spectral region are relatively insensitive
to the assumed isotopic fraction and the difference in derived Ba
abundances when the two assumptions on the isotopic fraction
are made is typically less than 0.02 dex. We have also updated
the overall log gf values of each Ba line to those recommended
by Fuhr & Wiese (2005).
Lanthanum. Lanthanum abundances are estimated from EWs
of La ii lines, taking into account the hfs for these lines. The
log gf values and the atomic data for the hfs were taken from
Ivans et al. (2006). The La ii lines for which hyperfine structure
data are not available in Ivans et al. (2006) are not used in theabundance derivation.
Neodymium and Samarium. The abundances of neodymium
and samarium have been obtained from EWs of Nd ii and Sm ii
lines, respectively. The log gf values for these lines have been
adopted from Ivans et al. (2006), in which the values from recent
laboratory measurements were employed.
Europium. Europium abundances have been determined fromthe spectral synthesis of Eu ii 4129.7 and/or 6645.1 Å lines.
An example of the fitted spectrum is shown in Figure 1. The
abundance fraction of two naturally occurring isotopes, 151 Eu
and 153 Eu, is assumed to be 151 Eu≡151 Eu/(151 Eu+153 Eu) = 0.5,
which is roughly consistent with the r-process component of thesolar-system meteorite abundance (Anders & Grevesse 1989).
Hyperfine and isotopic structures for these lines were calculated
based on the data listed in Lawler et al. (2001). The oscillator
strengths were also taken from Lawler et al. (2001). The relative
strengths of the transitions were computed with a standard
manner as described in Lawler et al. (2001). For lines of otherspecies within a few Å of each Eu ii line, we have adopted thelog gf values from the current version of the Kurucz (1995)
database. The Eu ii 6645 Å line is probably contaminated by a
Si i line at 6645.21 Å in the adopted line list. The log gf valueof this line was slightly modified so that the calculated synthetic
5

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba

Figure 1. Example spectra (circles) of the two sample stars for wavelength regions around the Eu II 4129.7 Å (left) and 6645.1 Å (right) lines. The adopted atmosphericparameters of these stars are indicated on the top of the left panels as “Teff (K)/log g/[Fe/H].” The solid lines show the best-fit synthetic spectra and the dotted lines
show the spectra for Δ log (Eu) = ±0.1 dex from the best-fit values.

with a scatter of 0.10 dex. The large difference is partly attributed
to the difference in the adopted microturbulent velocity (ξ ). As
shown in Paper I, the present study has adopted systematicallylower ξ than that in Nissen & Schuster (2011), which resultsin the larger Ba abundances. Another possible cause for thediscrepancy is the difference in the adopted damping constant,for which the Unsöld (1955) approximation to the Van der Waals
constant, enhanced by a factor of 2.2, was employed in this study.
Figure 3 shows comparisons of derived [Fe/H], log (Zn),log (Y), and log (Eu) abundances with those from Roederer
et al. (2010) and Simmerer et al. (2004) for the 21 stars studied in common. The abundance results from these studies andthis work are summarized in Table 4. For [Fe/H], log (Zn),and log (Y), our abundances are systematically higher, where
the means of the difference (scatter) are Δ[Fe/H](TW-R10) =0.12 (0.17), Δ log (Zn)(TW-R10) = 0.12 (0.12), and Δ log (Y)(TW-R10) = 0.10 (0.21) dex. These offsets could partly be attributed to the difference in the adopted Teff , which is higher
in this study by ∼60 K on average than those of Roederer
et al. (2010) taken from Simmerer et al. (2004). The mean ofthe difference is smaller for log (Eu) (−0.01 dex) but with alarger scatter of 0.28 dex. The larger scatter for log (Y) andlog (Eu) is partly due to a few stars for which large discrepancy is found. For the sample star G 63−46 with Δ log (Y)(TW-R10) = 0.44 dex, the difference is likely attributed to∼160 K difference in the adopted Teff . For another star HD
128279 with Δ log (Eu)(TW-R10) = 0.67 dex, only a single
Eu ii line, which is close to the detection limit, is used for theabundance estimate in the present study. Thus, the uncertainty
in the synthetic spectral fitting may be mainly responsible forthe discrepancy.

spectrum for the Sun best reproduces the observed solar spectrain the wavelength range surrounding this line. Mashonkina et al.
(2012) reported that non-LTE correction for the Eu ii 4129 Å line
ranges from 0.05 to 0.12 dex and is generally larger for lower
gravity stars (i.e., giants). In order to avoid systematic errors dueto the non-LTE effect, we present the results considering dwarfs
and giants separately when comparing the Eu abundancesamong the subsamples (see Section 4.1.8).
3.2.3. Abundance Errors

Errors in the abundances are computed by taking into accountline-to-line scatter in derived abundances and uncertainty inthe adopted atmospheric parameters as in Paper I. The line-toline scatter in the abundances from individual lines is typicallysmaller than 0.10 dex. Errors in the mean of the abundances dueto the scatter are calculated as the line-to-line scatter divided
by a square root of the number of lines used to compute themean. When only one line is used to estimate the abundances,
we assume that the error is equal to the line-to-line scatter in Fe ilines, which are typically the most numerous. Abundance errors
due to the uncertainty in the adopted Teff , log g, and ξ values are
examined by changing these parameters by ±100 K, ±0.3 dex,
and ±0.3 km s−1 , respectively, in the abundance estimates.
The final errors are obtained by summing these contributionsin quadrature and are listed in Table 2.
3.3. Comparison with Other Studies
Figure 2 shows the comparison of derived abundance ratioswith those from Nissen & Schuster (2010, 2011) for the ninestars analyzed in common (G 112−43, G 53−41, G 125−13,
G 20−15, G 176−53, G 188−21, HD 111980, HD 105004,
and HD 193901). The abundance results of the two studies
are also summarized in Table 3. For [Na/Fe], [Ni/Fe], and[Zn/Fe], the derived abundances show an excellent agreement
within 0.01 dex with the rms scatter for the difference of0.07 dex. The [Cr/Fe], [Mn/Fe], [Cu/Fe], and [Y/Fe] in thetwo studies marginally agree within the mean differences of at
most 0.09 dex. The derived [Ba/Fe] ratios tend to be larger in thiswork than in Nissen & Schuster (2011) by 0.28 dex on average

3.4. [X/Fe]–Teff Correlation
The sample stars in the present study have various Teff values
in the range 4000–6900 K. In this subsection, we examine the[X/Fe]–Teff correlation among the sample stars and examine theextent to which such a correlation might affect the abundancecomparison between the thick disk, inner halo, and outer halosubsamples.
6

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba

Figure 2. Comparison of the derived [Fe/H], [Na/Fe], [Cr/Fe], [Mn/Fe], [Ni/Fe], [Cu/Fe], [Zn/Fe], [Y/Fe], and [Ba/Fe] abundance ratios with those from Nissen& Schuster (2010, 2011).
Table 3Comparison with Nissen & Schuster (2010, 2011)
Star Name

NS10/TW

Teff
(K)

log g
(dex)

ξ
(km s−1 )[Fe/H](dex)[Na/Fe](dex)[Cr i/Fe](dex)[Mn/Fe](dex)[Ni/Fe](dex)[Cu/Fe](dex)[Zn/Fe](dex)[Y/Fe](dex)[Ba/Fe](dex)

Classification

G112−43

TW
NS
TW
NS
TW
NS
TW
NS
TW
NS
TW
NS
TW
NS
TW
NS
TW
NS

6176
6074
6070
5859
6079
5848
5798
5778
6042
6027
6115
5754
5753
5523
5908
5650
6170
5974

4.0
4.0
4.6
4.3
4.8
4.3
4.0
4.0
4.3
4.3
5.0
4.3
5.0
4.5
4.9
4.4
4.5
4.2

1.4
1.3
0.8
1.3
0.8
1.5
1.2
1.5
1.2
1.6
0.4
1.2
0.2
1.0
0.3
1.2
1.1
1.5

−1.33
−1.25
−1.15
−1.20
−1.35
−1.43
−1.13
−1.08
−1.62
−1.49
−0.60
−0.82
−1.26
−1.34
−0.95
−1.09
−1.29
−1.32

−0.11
−0.11
0.10
0.23
−0.20
−0.17
0.12
0.03
···
−0.18
−0.09
−0.05
−0.30
−0.36
−0.28
−0.27
0.02
−0.04

−0.04
0.00
−0.09
−0.03
−0.14
−0.09
−0.14
−0.02
−0.17
−0.03
−0.08
−0.03
−0.06
−0.01
−0.05
−0.04
−0.09
−0.04

−0.20
−0.19
−0.32
−0.38
−0.32
−0.39
−0.30
−0.34
−0.26
−0.32
−0.12
−0.21
−0.31
−0.35
−0.18
−0.34
−0.27
−0.34

0.00
−0.02
−0.14
−0.09
−0.07
−0.09
0.00
0.00
−0.01
−0.05
−0.12
−0.06
−0.08
−0.12
−0.13
−0.14
−0.02
−0.01

−0.49
−0.52
−0.75
···
−0.80
···
···
−0.32
···
···
···
−0.22
···
−0.57
−0.55
−0.62
···
−0.41

0.30
0.30
0.04
0.03
0.03
0.05
0.15
0.15
0.11
0.11
0.01
0.04
0.04
0.08
−0.04
−0.03
0.11
0.15

−0.16
−0.14
0.10
0.05
−0.10
−0.22
0.23
0.16
−0.16
−0.14
0.02
−0.15
−0.05
−0.26
−0.04
−0.20
0.26
0.15

−0.21
−0.27
0.54
0.24
−0.01
−0.31
0.32
0.07
0.02
−0.22
0.21
−0.16
0.12
−0.26
0.15
−0.19
0.32
0.03

OH
Low-α
IH
Low-α
IH
(High-α)OH
High-αOH(Low-α)
IH
Low-α
OH
Low-α
IH
Low-α
IH/TD
High-αG53−41
G125−13
HD111980
G20−15
HD105004
G176−53
HD193901
G188−22

Figures 4–6 show the [X/Fe] plotted against Teff . The left
and right rows of each figure show the plots for [Fe/H]  −2and <−2, respectively. The three sizes of symbols represent thethree metallicity intervals with the larger symbols corresponding
to higher metallicities (see the caption of Figure 4). A slope ofthe linear regression line, which is calculated by a two-sigma
clipping algorithm, is indicated in the top of each panel.
It can be seen from the figures that some of the elements
show a slope larger than 3σ in the [X/Fe]–Teff plane in one
or both metallicity range(s). In [Fe/H]  −2 (the left columns
7The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba

Figure 3. Comparison of the derived [Fe/H], log (Zn), log (Y), and log (Eu) abundances with those from Roederer et al. (2010).
Table 4Comparison with Roederer et al. (2010)
Name

Teff,TW Teff,R10 log gTW log gR10 ξTW ξR10 [Fe/H] TW [Fe/H] TW log (Zn)TW log (Zn)R10 log (Y)TW log (Y)R10 log (Eu)TW log (Eu)R10
(K)
(dex)
(km s−1 )
(dex)
(dex)
(dex)
(dex)

HD 107752
HD 119516
HD 124358
HD 128279
HD 85773
G17−25
HD 214362
HD 218857
G153−21G176−53
G188−22
G63−46
G23−14
HD 105546
HD 108317
HD 122956
HD 171496
HD 184266
HD 188510
HD 193901
HD 210295

4826
5605
4745
5328
4370
5174
5783
5107
5566
5753
6170
5867
5061
5179
5284
4609
4795
5618
5654
5908
4763

4649
5382
4688
5290
4268
4966
5727
5103
5700
5593
5827
5705
5025
5190
5234
4508
4952
6000
5564
5750
4750

1.6
1.9
1.4
3.2
0.7
4.8
2.0
2.7
3.9
5.0
4.5
4.6
3.1
2.3
2.9
1.6
1.9
1.6
5.0
4.9
2.2

1.6
2.5
1.6
3.0
0.5
4.3
2.6
2.4
4.4
4.5
4.3
4.2
3.0
2.5
2.7
1.6
2.4
2.7
4.5
4.5
2.5

1.9
1.9
1.8
1.4
2.0
0.0
2.3
1.7
0.9
0.2
1.1
0.6
1.1
1.8
1.5
1.7
1.1
2.4
0.1
0.3
1.3

2.0
2.5
2.1
1.5
2.0
0.8
2.0
1.9
1.4
1.2
1.2
1.3
1.3
1.6
2.0
1.6
1.4
3.0
1.0
1.5
1.6

−2.78
−1.92
−1.70
−2.18
−2.44
−1.14
−1.90
−1.91
−0.65
−1.26
−1.29
−0.62
−1.47
−1.44
−2.27
−1.71
−0.64
−1.68
−1.47
−0.95
−1.24

−2.78
−2.26
−1.91
−2.51
−2.62
−1.54
−1.87
−1.90
−0.70
−1.34
−1.52
−0.90
−1.64
−1.48
−2.18
−1.95
−0.67
−1.43
−1.32
−1.08
−1.46

1.92
2.67
2.73
2.45
2.65
3.68
2.76
2.75
4.16
3.38
3.41
4.12
3.19
3.24
2.43
2.92
4.23
2.99
3.15
3.59
3.52

1.93
2.32
2.64
2.15
2.56
3.40
2.71
2.64
4.06
3.18
3.24
3.86
3.05
3.29
2.40
2.87
4.11
3.19
3.01
3.36
3.37

−0.84
0.02
−0.13
···
−0.80
···
0.22
−0.09
1.36
0.90
1.18
1.70
0.61
0.76
−0.22
0.29
1.22
0.22
0.71
1.21
1.00

−0.90
−0.43
−0.22
−1.04
−0.93
0.87
0.32
−0.19
1.46
0.63
0.94
1.26
0.46
0.74
−0.39
0.16
1.40
0.69
0.44
0.83
0.85

−1.90
−0.98
−1.02
−1.60
−1.95
−0.33
−1.17
−1.55
0.01
−0.21
−0.44
0.20
−0.52
−0.63
−1.22
−0.83
−0.23
−1.08
−0.55
0.06
−0.46

−1.99
−1.43
−0.94
−2.27
−1.84
0.00
−0.82
−1.42
0.34
−0.32
−0.60
−0.05
−0.58
−0.56
−1.32
−0.79
0.11
−0.43
−0.52
−0.10
−0.34

exceed 0.18 dex/1000 K, which may affect the abundancecomparison between the subsamples. In [Fe/H] < −2 (rightcolumns of Figures 4–6), the abundance ratios of Sc, V, Cr i,
Y, Zr, Nd, and Sm show a slope of >3σ with Teff . Particularlyof Figures 4–6), a significant [X/Fe]–Teff correlation can berecognized for V, Cr i, Co, Nd, Sm, and Eu. For V, Cr, andCo, the slopes are 0.13 dex/1000 K, which is comparable tothe observational errors, while for Nd, Sm, and Eu, the slopes
8The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba

Figure 4. Abundance ratios of Na, Sc, V, Cr i, Cr ii, and Mn plotted against the adopted Teff values for the sample stars with the metallicity [Fe/H]  −2 (left)and [Fe/H] < −2 (right). The size of the symbols corresponds to metallicity; in the left (right) panel, small: −2.0  [Fe/H] < −1.5 ([Fe/H] < −3.0), medium:−1.5  [Fe/H] < −1.0 (−3.0  [Fe/H] < −2.5), and large: −1.0  [Fe/H] (−2.5  [Fe/H] < −2.0). A dotted line in each panel shows the result of a least-squares
fit to a straight line [X/Fe] = b + aTeff . The slope a of the fit is indicated in each panel.

may be different among different species and lines used in
the analysis, we do not correct [X/Fe] values to vanish the[X/Fe]–Teff slopes in the following analysis. Instead, for theelements with the large slopes, we separately treat dwarfs
and giants in the abundance comparisons among the threesubsamples.

large slopes for [Nd/Fe] and [Sm/Fe] versus Teff plots are partly
attributed to paucity of data points in the range Teff > 5000 K,
in which Nd and Sm abundances are below detection limits formany of the sample stars. Since the upper limits for some ofthe sample stars in this temperature range are below [Nd/Fe],[Sm/Fe] ∼ 0.5 dex, the apparent extreme slopes are likely
artificial.
The abundance–Teff correlations as indicated above or thediscrepancy in the derived abundances between dwarf and giantstars have been reported in previous studies. Bonifacio et al.
(2009) compare abundances in the sample of dwarfs and giantsin the metallicity range of −4  [Fe/H]  −2. They reported
that the dwarf versus giant discrepancy presents for many
elements. In particular, the Sc, Cr, Mn, Zn, and Co abundanceswere found to be higher in dwarf stars than in giant stars. Thiseffect is also seen in our sample for Sc and Cr i in a similar
metallicity range. On the other hand, our sample does not show
a significant dwarf/giant discrepancy for Mn, Zn, and Co.
The exact reason for the discrepancy is currently unclear.
Bonifacio et al. (2009) suspect that granulation in the stellar
atmospheres (three-dimensional effects) and/or departure from
LTE might be responsible for the discrepancy observed in someelements. Since the magnitudes and direction of these effects

4. RESULTS
4.1. Distribution of the Sample Stars in [X/Fe]–[Fe/H] Planes
Figures 7 and 8 show the abundance ratios ([X/Fe]) plotted against [Fe/H] for the thick disk stars (crosses), theinner halo stars (filled circles), the outer halo stars (filled
triangles), and their intermediate populations (thick disk/
inner halo: open circles; inner halo/outer halo: open triangles). For particularly interesting elements, namely, Na, Ni,
Zn, and Eu, we additionally discuss the behavior of low[Mg/Fe] stars with [Mg/Fe] < 0.1, which are analogous to
the low-α stars in Nissen & Schuster (2010) for the purpose of examining the consistency of the present resultswith those of Nissen & Schuster (2010). The low-[Mg/Fe]
stars are marked with gray circles in the corresponding panels in Figures 7 and 11. Table 5 summarizes the means (μ) and9[Fe/H] > −1.5

[X/Fe]
μa
Na

Sc

V

Cr i

Cr ii

10
Mn

Co

Ni

Cu

Zn

Sr

[Fe/H]  −2.5Ncμd (N)dμg (N)eμσNμd (N)μg (N)

0.11
0.17
0.19
0.09
0.07
0.06
0.19
0.13
0.12
0.11
0.12
0.07
0.06
0.05
0.05
0.18
0.08
0.08
0.12
0.15
0.10
0.05
0.05
0.07
0.23
0.29
0.18
0.08
0.09
0.12
···
−0.07 0.10
0.02 0.07

8
16
11
8
18
11
8
15
10
8
18
11
8
18
11
8
18
11
5
5
5
8
18
11
7
11
6
8
18
11

0.16 ± 0.06 (2)−0.14 ± 0.02 (11)−0.23 ± 0.03 (6)0.29 ± 0.07 (2)0.15 ± 0.02 (12)0.16 ± 0.01 (6)0.22 ± 0.11 (2)0.12 ± 0.04 (9)0.15 ± 0.02 (6)−0.07 ± 0.00 (2)−0.09 ± 0.01 (12)−0.08 ± 0.01 (6)0.22 ± 0.01 (2)0.17 ± 0.01 (12)0.17 ± 0.02 (6)−0.03 ± 0.10 (2)−0.27 ± 0.02 (12)−0.28 ± 0.02 (6)−0.07 ± 0.13 (1)0.20 ± 0.03 (2)0.14 ± 0.02 (4)0.02 ± 0.02 (2)−0.09 ± 0.01 (12)−0.08 ± 0.03 (6)−0.07 ± 0.01 (2)−0.53 ± 0.09 (6)−0.49 ± 0.12 (1)0.22 ± 0.04 (2)0.03 ± 0.02 (12)0.07 ± 0.05 (6)...
−0.07 ± 0.04 (6)0.02 ± 0.03 (4)0.13 ± 0.04 (4)−0.13 ± 0.07 (3)−0.30 ± 0.14 (1)0.22 ± 0.02 (4)0.15 ± 0.06 (3)−0.03 ± 0.12 (1)−0.10 ± 0.08 (4)−0.11 ± 0.03 (3)−0.12 ± 0.17 (1)−0.14 ± 0.07 (4)−0.29 ± 0.11 (3)−0.26 ± 0.18 (1)0.15 ± 0.03 (4)0.13 ± 0.05 (3)0.09 ± 0.09 (1)0.01 ± 0.07 (4)−0.22 ± 0.01 (3)−0.17 ± 0.08 (1)−0.04 ± 0.07 (4)0.04 ± 0.12 (2)−0.05 ± 0.16 (1)0.00 ± 0.01 (4)−0.08 ± 0.01 (3)−0.17 ± 0.09 (1)−0.25 ± 0.05 (4)−0.20 ± 0.27 (2)−0.85 ± 0.12 (1)0.22 ± 0.04 (4)0.14 ± 0.06 (3)−0.07 ± 0.19 (1)
...
...
...

−0.08
−0.15
−0.28
0.15
0.13
0.05
0.04
0.14
−0.05
−0.23
−0.14
−0.17
0.18
0.19
0.15
−0.33
−0.34
−0.37
0.12
0.11
0.06
−0.06
−0.08
−0.10
−0.57
−0.70
−0.74
0.14
0.13
0.01
−0.03
0.13
−0.09

0.08
0.14
0.18
0.05
0.09
0.10
0.02
0.24
0.13
0.07
0.08
0.09
0.06
0.05
0.06
0.08
0.10
0.10
0.19
0.13
0.13
0.05
0.06
0.09
0.12
0.24
0.27
0.00
0.15
0.10
0.23
0.15
0.25

2
8
5
2
13
20
2
8
11
2
13
20
2
13
19
2
13
20
2
8
9
2
12
19
1
4
8
2
12
19
1
5
6

...
−0.19 ± 0.05 (4)−0.01 ± 0.09 (1)...
0.18 ± 0.04 (6)0.08 ± 0.05 (5)...
0.28 ± 0.08 (5)−0.19 ± 0.13 (1)...
−0.07 ± 0.01 (6)−0.12 ± 0.03 (5)...
0.20 ± 0.02 (6)0.15 ± 0.02 (5)...
−0.35 ± 0.02 (6)−0.34 ± 0.03 (5)...
0.13 ± 0.03 (5)0.10 ± 0.01 (4)...
−0.08 ± 0.04 (5)−0.14 ± 0.04 (5)...
−0.45 ± 0.12 (1)...
...
0.04 ± 0.04 (5)0.01 ± 0.06 (5)...
0.13 ± 0.07 (5)−0.04 ± 0.11 (5)−0.03 ± 0.14 (1)−0.11 ± 0.04 (2)−0.35 ± 0.02 (4)0.19 ± 0.08 (1)0.07 ± 0.02 (4)0.03 ± 0.02 (8)0.05 ± 0.13 (1)−0.09 ± 0.05 (3)−0.09 ± 0.02 (6)−0.18 ± 0.07 (1)−0.22 ± 0.02 (4)−0.26 ± 0.02 (8)0.22 ± 0.09 (1)0.17 ± 0.02 (4)0.13 ± 0.02 (8)−0.27 ± 0.07 (1)−0.23 ± 0.04 (4)−0.42 ± 0.03 (8)−0.02 ± 0.14 (1)0.09 ± 0.13 (3)0.32 ± 0.22 (1)−0.10 ± 0.10 (1)−0.07 ± 0.03 (4)−0.13 ± 0.02 (8)−0.57 ± 0.12 (1)−0.67 ± 0.08 (2)−0.87 ± 0.07 (6)0.14 ± 0.18 (1)0.23 ± 0.10 (4)0.01 ± 0.04 (8)
...
...
...

0.10
−0.13
−0.28
0.21
0.14
0.14
0.02
0.06
0.12
−0.13
−0.14
−0.10
0.18
0.16
0.17
−0.08
−0.25
−0.27
−0.05
0.08
0.10
−0.01
−0.08
−0.12
−0.26
−0.49
−0.54
0.20
0.06
0.03

6
4

μ

σ

···
···
−0.06 0.14
0.17 0.10
0.14 0.26
0.24 0.26···
0.17 0.17−0.27 0.16
−0.06 0.07
−0.06 0.10
−0.23 0.12
0.18 0.10
0.16 0.07
0.08 0.06−0.68 0.07
−0.44 0.12
−0.34 0.16
0.05 0.10
0.17 0.03······−0.06 0.04
0.03 0.12······−0.68 0.12···
−0.01 0.27
0.11 0.10···
0.09 0.17
0.06 0.12

N

1
1
3
6
1
1
1
3
5
1
1
1
1
2
1
1
2

2
2

1
2
2
3
4

μd (N)μg (N)

...
...
...
...
0.28 ± 0.04 (2)0.36 ± 0.11 (4)...
0.17 ± 0.17 (1)...
...
−0.00 ± 0.01 (2)−0.18 ± 0.08 (3)
...
...
...
...
−0.35 ± 0.16 (1)...
...
0.15 ± 0.16 (1)...
...
−0.09 ± 0.18 (1)0.12 ± 0.08 (1)
...
...
...
...
−0.20 ± 0.21 (1)...
...
0.16 ± 0.12 (2)0.08 ± 0.08 (3)...
...
−0.06 ± 0.14 (1)...
−0.16 ± 0.08 (1)0.04 ± 0.09 (1)...
...
−0.27 ± 0.16 (1)...
−0.17 ± 0.14 (1)−0.31 ± 0.08 (1)...
0.16 ± 0.07 (1)0.08 ± 0.06 (1)...
−0.53 ± 0.14 (1)−0.34 ± 0.16 (1)...
0.19 ± 0.17 (1)...
...
−0.04 ± 0.11 (1)−0.06 ± 0.08 (1)...
...
−0.68 ± 0.12 (1)...
0.18 ± 0.13 (1)0.04 ± 0.21 (1)...
−0.04 ± 0.18 (1)...

Ishigaki, Aoki, & Chiba

TD
IH
OH
TD
IH
OH
TD
IH
OH
TD
IH
OH
TD
IH
OH
TD
IH
OH
TD
IH
OH
TD
IH
OH
TD
IH
OH
TD
IH
OH
TD
IH
OH

−2.5 < [Fe/H]  −1.5

σb

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Table 5Means and Standard Deviations in the Abundance Ratios

[Fe/H] > −1.5

[X/Fe]

Y

Zr

Ba

La

11
Nd

Sm

Eu

TD
IH
OH
TD
IH
OH
TD
IH
OH
TD
IH
OH
TD
IH
OH
TD
IH
OH
TD
IH
OH

−2.5 < [Fe/H]  −1.5

[Fe/H]  −2.5

μa

σb

Nc

μd (N)dμg (N)eμσNμd (N)μg (N)μσNμd (N)μg (N)

−0.11
0.06
−0.02
0.12
0.29
0.34
0.09
0.21
0.09
0.14
0.42
0.46
0.27
0.57
0.50
0.19
0.50
0.60
0.19
0.39
0.50

0.14
0.20
0.12
0.21
0.18
0.12
0.07
0.29
0.15
0.08
0.48
0.35
0.19
0.30
0.17
0.11
0.34
0.16
0.16
0.27
0.21

8
17
9
4
4
5
8
18
10
5
9
6
8
18
10
8
10
7
8
18
10

−0.05 ± 0.16 (2)0.09 ± 0.06 (12)−0.00 ± 0.05 (6)0.10 ± 0.20 (2)0.40 ± 0.15 (2)0.32 ± 0.06 (4)0.05 ± 0.06 (2)0.27 ± 0.10 (12)0.11 ± 0.07 (6)···0.73 ± 0.29 (4)0.36 ± 0.12 (3)0.55 ± 0.15 (2)0.68 ± 0.07 (12)0.62 ± 0.07 (5)0.21 ± 0.10 (2)0.82 ± 0.16 (4)0.68 ± 0.13 (3)0.22 ± 0.09 (2)0.48 ± 0.04 (12)0.49 ± 0.08 (6)−0.16 ± 0.07 (4)0.07 ± 0.03 (3)···−0.04 ± 0.10 (1)······0.13 ± 0.03 (4)0.07 ± 0.07 (3)···0.14 ± 0.05 (4)0.15 ± 0.09 (3)0.25 ± 0.14 (1)0.16 ± 0.01 (4)0.23 ± 0.01 (3)0.35 ± 0.14 (1)0.14 ± 0.05 (4)0.21 ± 0.06 (3)0.47 ± 0.16 (1)0.09 ± 0.07 (4)0.01 ± 0.25 (3)0.23 ± 0.21 (1)

−0.10
−0.12
−0.23
0.25
0.20
−0.06
0.06
−0.03
−0.11
0.31
0.43
0.34
0.33
0.30
0.23
0.50
0.35
0.26
0.51
0.20
0.34

0.09
0.32
0.20
0.01
0.37
0.11
0.08
0.53
0.26
0.11
0.63
0.47
0.00
0.53
0.26
0.07
0.50
0.13
0.03
0.42
0.32

2
11
16
2
6
7
2
13
20
2
5
7
2
10
14
2
7
10
2
9
20

···
0.02 ± 0.10 (4)−0.14 ± 0.12 (4)···0.40 ± 0.24 (3)······0.00 ± 0.08 (6)0.03 ± 0.08 (5)······0.36 ± 0.14 (1)···0.62 ± 0.16 (3)0.57 ± 0.04 (3)
···
···
···
···
0.35 ± 0.12 (5)0.59 ± 0.06 (5)−0.04 ± 0.09 (1)−0.36 ± 0.09 (4)−0.32 ± 0.06 (8)0.25 ± 0.11 (1)0.01 ± 0.15 (3)−0.03 ± 0.04 (6)0.12 ± 0.22 (1)−0.34 ± 0.20 (4)−0.23 ± 0.09 (8)0.39 ± 0.15 (1)0.16 ± 0.08 (2)0.12 ± 0.09 (4)0.33 ± 0.10 (1)−0.12 ± 0.18 (4)0.09 ± 0.05 (7)0.45 ± 0.10 (1)0.11 ± 0.12 (4)0.25 ± 0.05 (7)0.48 ± 0.10 (1)−0.01 ± 0.23 (4)0.16 ± 0.11 (8)

−0.04
−0.07
−0.46
0.42
0.20
−0.16
−0.50
−0.25
−0.28

0.09
0.27
0.05
0.09
0.09
0.14
0.12
0.14
0.37

1
2
2
1
1
1
1
3
5

···
0.12 ± 0.12 (1)
···
···
···
···
···
−0.24 ± 0.14 (2)−0.03 ± 0.05 (3)
···
···
···
···
···
···
···
···
···
···
0.70 ± 0.20 (2)1.07 ± 0.31 (2)···−0.26 ± 0.09 (1)−0.49 ± 0.10 (1)···0.20 ± 0.09 (1)−0.16 ± 0.14 (1)···−0.25 ± 0.12 (1)−0.46 ± 0.17 (1)······−0.02 ± 0.14 (1)···0.17 ± 0.10 (1)−0.11 ± 0.11 (1)···0.33 ± 0.11 (1)0.03 ± 0.12 (1)···0.36 ± 0.21 (1)0.09 ± 0.11 (1)

−0.02
0.17
−0.11
0.33
0.03
0.06
0.43
0.09

···
···
0.14

1

0.10
0.11

1
1

0.11
0.12
0.21
0.10
0.11

1
1
1
2
1

···
···

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Table 5(Continued)

Notes.
a Means of the abundance ratios within a given [Fe/H] interval.
b Standard deviations of the means.
c The number of stars used to calculate the μ and σ .
d A mean and its error for dwarf stars only. The number of stars is given in parentheses.
e A mean and its error for giant stars only. The number of stars is given in parentheses.
Ishigaki, Aoki, & Chiba

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba

Figure 5. Same as Figure 4, but for Co, Ni, Cu, Zn, Sr, and Y.

in a given metallicity range with a possible exception of the twostars having [Na/Fe] > 0.0.
The observed trends of [Na/Fe] ratios with [Fe/H] for our
three subsamples are in agreement with previous studies. Reddy
et al. (2006) reported modestly enhanced [Na/Fe] ratios anddecreasing [Na/Fe] with decreasing [Fe/H] for their sample ofthick disk stars with [Fe/H] < −0.6, which is similar to thetrend seen in the plot in Figure 7. The non-LTE (re)analysis
of sodium abundances by Takeda et al. (2003) suggests that
[Na/Fe] ratios in their sample of thick disk stars are nearsolar, which is roughly consistent with the present result if thesuggested non-LTE correction up to ∼−1.0 dex is applied toour sample.
For the inner and outer halo stars, our results suggest that the[Na/Fe] ratios are likely correlated with both kinematics and[α/Fe], which qualitatively supports previous findings by
Stephens & Boesgaard (2002), Fulbright (2002), and Nissen& Schuster (2010, 2011). As an example, Fulbright (2002) reported that a fraction of Na-poor stars defined as [Na/Fe] <−0.36 in his sample increases for stars with a large apocentric
distance (Rapo > 20 kpc). The suggested dependence is similar
to the present result in that the outer halo subsample, which
includes stars with Rapo > 15 kpc, tends to show relatively low[Na/Fe] ratios than the inner halo stars. We will discuss thelikely [Na/Fe]–kinematics correlations in Section 4.2. Nissen& Schuster (2010) reported the distinct [Na/Fe] ratios for the

scatters (σ ) of the abundance ratios for the three subsamples(in the second column, “TD,” “IH,” and “OH,”) for the thick
disk, inner halo, and outer halo subsamples, respectively. This
table also includes the means and scatters of the abundance ratios taking into account dwarfs (μd ) or giants (μg ) alone. The
anomalous star, LP 751−19 (see Section 3.2.1), is excluded in
the calculation of the means and scatters.
4.1.1. Sodium, Scandium, and Vanadium

The thick disk, inner halo, and outer halo subsamples show
different trends and scatters in the [Na/Fe]–[Fe/H] diagram.
First, the thick disk stars are modestly enhanced in [Na/Fe]
ratios at metallicities [Fe/H] > −1.5 with a mean abundanceof [Na/Fe] = 0.10 dex, while most of the inner and outer halostars show lower ratios at similar metallicities. Second, scatterin the [Na/Fe] ratios in this [Fe/H] range is relatively small(0.11 dex) for the thick disk stars, which is comparable to theobservational error, while the inner and outer halo stars show
larger scatter (0.17 dex). Third, the mean [Na/Fe] ratio forthe outer halo stars (−0.28 ± 0.06 dex) is lower than that of theinner halo stars (−0.13 ± 0.04 dex). There are also abundance
differences among the three subsamples when only dwarfs or
giants are taken into account (see Table 5). Furthermore, among
the inner halo and outer halo stars, those with lower [Mg/Fe]
ratios (gray circled symbols) tend to show lower [Na/Fe] ratios
12

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba

Figure 6. Same as Figure 4, but for Zr, Ba, La, Nd, Sm, and Eu.

low- and high-α stars in the metallicity range of −1.6 <[Fe/H] < −0.4, for which a similar trend is apparent in
Figure 7. To summarize, the inner and outer halo stars, at leastin part, show lower [Na/Fe] ratios than the thick disk stars andthe lowest [Na/Fe] stars in our sample tend to have outer halo
kinematics and/or low [Mg/Fe] ratios.
The observed difference in the [Na/Fe] ratios among thethree subsamples suggests that progenitors of the thick disk
and the inner and outer halos experienced largely different
chemical enrichment histories. Sodium is mainly synthesized
during hydrostatic carbon burning in massive stars (Woosley& Weaver 1995) and its yield is known to be dependent on
the mass and metallicity of the progenitor star (Kobayashi
et al. 2006). Therefore, the IMF, typical metallicities, and/or
the relative contribution of Fe from Type Ia SNe may be
responsible for determining [Na/Fe] abundance ratios of theprogenitor systems. One possible interpretation for the originof the observed lower [Na/Fe] stars in our inner and outer halosubsamples is that higher mass stars were deficient in theirprogenitors compared to those of the thick disk. Alternatively,
metals ejected from massive stars more easily escape in theprogenitor of the inner and outer halos than in those of the
thick disk. Star formation rate is another factor that could affect[Na/Fe] ratios since it determines the relative contribution of
Na predominantly from Type II SNe to Fe from Type Ia SNe.
In order to examine which of the above factors is the mostimportant for explaining the observed [Na/Fe] difference among
the three subsamples, a chemical evolution modeling whichtakes into account the difference in star formation environment
during the formation of the three components is necessary.
The [Sc/Fe] ratios for the thick disk, inner halo, and outer
halo subsamples are all enhanced in the metallicity range[Fe/H] > −1.5 with a modest decreasing trend toward lower[Fe/H] with small scatter (<0.10 dex). A significant differencein this trend among the three subsamples is not found. In[Fe/H] below ∼−1.5, the trend appears to be flattened at[Sc/Fe] ∼ 0.10 dex showing a much larger scatter (>0.25 dex;
Table 5). The observed enhanced [Sc/Fe] ratios are generally
consistent with those found in previous studies for stars with
[Fe/H] > −1.5 (Prochaska et al. 2000; Reddy et al. 2006; Zhao
& Magain 1990; Nissen et al. 2000; Cayrel et al. 2004).
The [V/Fe] ratios for the three subsamples are nearly flat at
the solar value in the metallicity below [Fe/H] ∼ −1.0 withouta significant difference among the subsamples. The near-solar
value for the halo stars is in agreement with previous studies(Gratton & Sneden 1991; Lai et al. 2008). The enhancement in
the [V/Fe] ratio in [Fe/H] > −1.0 is reported by Prochaska
et al. (2000) and Reddy et al. (2006) for the thick disk stars,
while a modest enhancement (0.02–0.12 dex) is also seen in our
sample.
In the theoretical calculation of Woosley & Weaver (1995),
both Sc and V are produced through explosive oxygen, silicon,
13

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba

Figure 7. Abundance ratios for Na, Sc, V, Cr i, Cr ii, Mn, Co, Ni, Cu, and Zn plotted against [Fe/H]. The crosses, filled circles, and filled triangles indicate the sample
stars with PTD > 0.9 (the thick disk stars), PIH > 0.9 (the inner halo stars), and POH > 0.9 (the outer halo stars), respectively. The open circles show the stars whose
kinematics are intermediate between the thick disk and the inner halo (PTD , PIH  0.9 and PTD , PIH  POH ), while open triangles indicate the stars whose kinematics
are intermediate between the inner and the outer halo (PIH , POH  0.9 and PIH , POH  PTD ). Symbols marked with a gray open circle in the [Na/Fe], [Ni/Fe], and[Zn/Fe] panels represent the sample stars with [Mg/Fe] < 0.1.
(A color version of this figure is available in the online journal.)

As far as the [Cr ii/Fe] ratio is concerned, which is presumably less affected by the overionization, the abundance ratios
show a negligible scatter of 0.06 dex at most for all of the
thick disk, inner halo, and outer halo subsamples over the whole[Fe/H] range. This result is consistent with that of a more precise analysis of Cayrel et al. (2004), which reported the 0.05 dex
scatter over the metallicity −4 < [Fe/H] < −2. The present results further confirm the small cosmic scatter in the [Cr/Fe]
ratios in the solar neighborhood stars independent of theirkinematics and metallicity.

and neon burning in massive stars and thus their yield issensitive to various parameters of SN explosions. The similarity
in the observed trend in [Sc/Fe] and [V/Fe] among the threesubsamples may indicate that the astrophysical sites for the
production of these elements are largely common among thedifferent Galactic populations.
4.1.2. Chromium

For chromium, we first note that the abundances derived
from the neutral and ionized species ([Cr i/Fe] and [Cr ii/Fe],
respectively) are systematically different, as can be seen in theplot in Figure 7. The [Cr i/Fe] ratios are subsolar for the whole[Fe/H] range and slightly decrease toward lower metallicity,while the [Cr ii/Fe] ratios are supersolar without any trends
with [Fe/H]. The observed discrepancy in the Cr abundance
ratios between those derived from neutral and ionized species
has previously been reported in the literature (Gratton & Sneden
1991; Lai et al. 2008; Bonifacio et al. 2009). The reason for thisdiscrepancy and which species is the more reliable indicator of
the true Cr abundance are not clear. Gratton & Sneden (1991)
suggested that the neutral species is affected by overionizationand thus use of Cr i lines would underestimate the overall Cr
abundance.

4.1.3. Manganese, Nickel, and Zinc

For all of the three subsamples, the [Mn/Fe] ratios show
an increasing trend with increasing [Fe/H] in the range[Fe/H] > −2. In a more metal-poor range, the trend appears to
be flattened with larger scatter. The increasing [Mn/Fe] trend
is consistent with previous studies for the thick disk and halostars in the range −1.5 < [Fe/H] < 0.0 (Prochaska et al. 2000;
Nissen et al. 2000; Reddy et al. 2006). Systematic differences
between the thick disk, inner halo, and outer halo stars are notclearly seen.
Manganese is thought to be mostly produced by explosive silicon burning in massive stars in their outer incomplete
14

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba

Figure 8. Same as Figure 7, but for neutron-capture elements (Sr, Y, Zr, Ba, La, Nd, Sm, and Eu). The two sample stars, G 18–24 and BD + 04◦ 2466, which are
identified as s-process-rich stars, are marked with larger circles.
(A color version of this figure is available in the online journal.)

show relatively small scatter (0.09 dex) of [Ni/Fe] ratios
about the mean value. Nissen & Schuster (2010) reported theunderabundance of Ni in low-α stars. Although the present study
does not have very high precision to convincingly confirm thisargument, the low-[Mg/Fe] stars (gray circled symbols) tend
to show lower [Ni/Fe] than the other stars (predominantly the
thick disk stars and thick disk/inner halo intermediate stars) in[Fe/H]  −1.2.
Nickel isotopes are produced in both deep layers of massive
stars and in Type Ia SNe (Woosley & Weaver 1995; Timmes et al.
1995). The observed difference in [Ni/Fe] among the thick disk
and the inner/outer halo stars at [Fe/H] > −2.5 may indicatethat the relative contribution of massive stars and Type Ia SNeto the Ni production is different among the progenitors of thesesubsamples.
The [Zn/Fe] ratios show an interesting difference between
the subsamples as can be seen in the bottom right panel of
Figure 7. The thick disk stars show supersolar [Zn/Fe] values
with relatively small scatter over all metallicities. The inner
halo stars show slightly lower [Zn/Fe] particularly in [Fe/H] >−1.0 with a modest scatter. The outer halo stars show lower[Zn/Fe] than the other two subsamples in the intermediate
metallicity range with a scatter similar to the inner halo stars.
One outer halo star G 112−43, which constitutes a commonproper-motion system with G 112−44, shows an exceptionallyhigh-[Zn/Fe] abundance compared to the other outer halo stars.
Such a peculiar abundance for this star is previously noted byNissen & Schuster (2010, 2011).
Nissen & Schuster (2011) reported that in −1.6 < [Fe/H] <−0.4, the thick disk stars and the high-α halo stars show
constantly high-[Zn/Fe] ratios while those of the low-α stars
show a mildly decreasing trend with [Fe/H]. In this metallicity,

Si-burning layers (Umeda & Nomoto 2005) and in Type Ia SNe(Iwamoto et al. 1999). In metallicities below [Fe/H] ∼ −1,
the low [Mn/Fe] ratios are mainly determined by the yields
from Type II SNe of massive stars (Tsujimoto & Shigeyama
1998). In [Fe/H]  −1, the increasing [Mn/Fe] with increasing[Fe/H] is interpreted as an onset of contribution fromType Ia SNe (Kobayashi et al. 2006). It was also suggested
that the dependence of Mn yields of Type Ia SNe on metallicity may contribute to the increase of [Mn/Fe] ratios with
[Fe/H] (Cescutti et al. 2008). The observed increase in the[Mn/Fe] ratios for the three subsamples, therefore, may indicate that Type Ia SNe have played some role for chemical
evolution in the progenitors of the thick disk, inner halo, andouter halo components. This interpretation favors the idea thatthe formation timescales for these progenitors were modestly
longer than those of the Type Ia SNe. The timescale for thechemical enrichment via Type Ia SNe is poorly constrained dueto the uncertainties in physical mechanisms that give rise to the
explosion (e.g., Maoz et al. 2010). Analyses of SN rates in galaxies and galaxy clusters suggest that the SN rate as a functionof the delay time from major star formation to SN explosions
is higher within a few Gyr and decreases toward longer timedelays (Totani et al. 2008; Maoz et al. 2010). Such estimates for
the SN delay time distributions as well as constraints on SN Mn
yields are essential to interpret the [Mn/Fe] ratios for each of
the three subsamples.
For the [Ni/Fe] ratios, a modest difference between the thickdisk and inner/outer halo subsamples can be recognized. As
indicated in Table 5, the thick disk stars show near-solar [Ni/Fe]
ratios while the inner halo and outer halo stars show ∼0.10 dex
lower [Ni/Fe] ratios in [Fe/H] > −2.5. The difference between
the two halo subsamples is not very clear. All of the subsamples
15

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba

a similar difference can be recognized in our sample: the[Zn/Fe] ratios of low-[Mg/Fe] stars (gray circled symbols) arelower than other stars on average. Additionally, the difference
seems to continue toward a lower metallicity of [Fe/H] ∼ −2.0
in the present sample. In the much lower metallicity range([Fe/H] < −2), both the inner halo and the outer halo stars showsupersolar [Zn/Fe] ratios in agreement with previous studies(Cayrel et al. 2004; Takeda et al. 2005; Lai et al. 2008).
At low metallicities ([Fe/H]  −2.5), Zn is thought to be
produced in the deep complete Si-burning region in massive
Type II SNe (Umeda & Nomoto 2005; Kobayashi et al. 2006).
Umeda & Nomoto (2005) suggest that the ejection of Zn is moreenhanced in higher energy SNe, which is called hypernovae.
Zinc is also thought to be synthesized through a weak s-processcomponent in massive stars (Timmes et al. 1995). The observed
supersolar [Zn/Fe] in our sample, therefore, could be a signatureof the hypernovae and/or neutron-capture process in massive
stars. At higher metallicities, the contribution from the sprocess in low-mass asymptotic giant branch (AGB) stars andFe production through Type Ia SNe could affect the observed[Zn/Fe] (Timmes et al. 1995). The [Zn/Fe] differences among
the three subsamples predominantly seen in [Fe/H] > −2.0
indicate different contributions of Type Ia and AGB products inthe progenitor interstellar medium (ISM) of these populations.
Additionally, the difference in mass and/or metallicity ofstars responsible for the Zn production may also affect theobserved abundance differences. Deeper understanding of the[Zn/Fe] among different populations requires more robust
estimates of Zn yields and their production timescales in various
astrophysical sites.

thick disk and the inner halo subsamples, while the trend is notclear in the outer halo subsample.
A large fraction of Cu is thought to be produced in a neutroncapture process in massive stars during their convective coreHe-burning and the shell C-burning (e.g., Pignatari et al. 2010).
The neutron source in these stars is mostly provided by the22
Ne(α,n)25 Mg reaction, where 22 Ne is produced from CNO
isotopes. Thus, in this scenario, the production of Cu depends
on the initial CNO composition of the progenitor star. This
production channel of Cu may explain the modest increasing[Cu/Fe] trend with [Fe/H] for the thick disk and the inner halostars. If this production channel is dominant at [Fe/H] > −1.5,
the lower [Cu/Fe] ratios for some inner/outer halo stars may
require extra enrichment of Fe presumably from Type Ia SNe.
4.1.5. Light Neutron-capture Elements: Sr, Y, and Zr

The [Sr/Fe] abundances are near-solar with scatter
0.07–0.25 dex for both of the inner halo and the outer halo subsamples (the top panel of Figure 8). Although only one thick disk
star has the Sr abundance measurement, the abundance of thisstar seems to agree with the inner/outer halo stars with similar
metallicity. The trends for the inner halo and the outer halo stars
are indistinguishable. The near-solar [Sr/Fe] in [Fe/H] > −2.0
is in agreement with previous studies (Gratton & Sneden 1994;
Burris et al. 2000), and the large dispersion in lower metallicity
is also consistent with previous studies (McWilliam et al. 1995;
Burris et al. 2000; Honda et al. 2004; François et al. 2007),
although the sample size with [Fe/H] < −2.0 in this study is
not large enough to quantify the scatter.
For [Y/Fe] ratios, the three subsamples show relatively large
abundance scatter in all metallicities. We note that two stars,
BD + 04◦ 2466 and G 18–24, with exceptionally high [Y/Fe],are the Y- and Ba-rich stars reported previously (Burris et al.
2000; Zhang et al. 2009; Ishigaki et al. 2010) and marked in
Figure 8 with large open circles. The former was found to bea spectroscopic binary (Jorissen et al. 2005), while the binary
nature of the latter is unclear (Latham et al. 2002).
The different behavior in the thick disk, inner halo, and outer
halo subsamples is modestly apparent in [Fe/H] > −1.5. The
thick disk stars show the [Y/Fe] slightly below the solar value
with a scatter of 0.14 dex. On the other hand, some of the inner
halo stars show the [Y/Fe] above the solar value. Finally, theouter halo stars show, again, near to subsolar [Y/Fe]. In thelower metallicity, the inner halo and the outer halo subsamples
both seem to show decreasing [Y/Fe] ratios with decreasing[Fe/H] with a large scatter of 0.20 dex.
Our results for the thick disk stars are generally consistent
with the previous studies of Bensby et al. (2005), Prochaska
et al. (2000), and Reddy et al. (2006). The decreasing [Y/Fe] forthe halo stars with decreasing [Fe/H] is also consistent with theresults of Gratton & Sneden (1994) and François et al. (2007).
[Zr/Fe] ratios in the range [Fe/H] > −1.5 show a hint ofdifference among the subsamples in that the inner/outer halostars seem to be more enhanced with [Zr/Fe] than the thickdisk stars. In the lower metallicity, [Zr/Fe] is higher in the thickdisk stars than in the halo stars. The [Zr/Fe] trend with [Fe/H]
for the halo stars is in agreement with that reported by Gratton& Sneden (1994). Such a trend in [Zr/Fe] ratios toward lower[Fe/H] has also been reported by François et al. (2007). In thisstudy, we show that the decreasing trend toward lower [Fe/H]
is seen only for the inner/outer halo stars with [Fe/H] < −2.0,
while it is not seen for the thick disk stars.

4.1.4. Cobalt and Copper

The [Co/Fe] ratios of the thick disk, inner halo, and outer
halo stars all show similar behavior against [Fe/H], namely, the[Co/Fe] is close to the solar value in [Fe/H] > −2.0, while itis ∼0.2 dex in lower metallicities. The scatter of 0.10–0.19 dex
may partly be attributed to the observational error, since only
one line, which is particularly strong in giant stars, is used for theabundance estimate in most of the sample stars. The enhanced[Co/Fe] ratios in [Fe/H] < −2.0 generally agree with the resultsof Cayrel et al. (2004) and Lai et al. (2008).
Cobalt is produced in both Type II and Type Ia SNe (Timmes
et al. 1995; Kobayashi et al. 2006). A chemical evolution modelof Timmes et al. (1995) suggests that the trend in [Co/Fe]
with respect to [Fe/H] is determined by the dependence ofCo ejection on mass/metallicity of massive stars and by the
production of Fe through Type Ia SNe. The observed similarity
in [Co/Fe] among the thick disk, inner halo, and outer halosubsamples is in contrast to the [Zn/Fe], for which abundance
differences among the three subsamples are identified. This
result suggests that, unlike [Zn/Fe], [Co/Fe] ratios are relatively
insensitive to star formation/chemical enrichment histories invarious progenitor systems.
The [Cu/Fe] ratios are plotted against [Fe/H] in the bottom
right panel of Figure 7. The downward arrows are overlaid
for the sample stars for which only an upper limit has beenobtained. The thick disk stars show near to subsolar [Cu/Fe]
ratios in the range −0.5 to 0.0. Both of the inner halo and outer
halo stars show lower [Cu/Fe] (∼−0.8) at metallicities up to[Fe/H] ∼ −1.0 with much larger scatter. The lower [Cu/Fe]
for the halo stars than the thick disk stars is consistent with theresults of Mishenina et al. (2002) and Reddy et al. (2006). An
increasing [Cu/Fe] trend with [Fe/H] is modestly seen in the16The Astrophysical Journal, 771:67 (25pp), 2013 July 1

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In the solar metallicity, a large fraction of Sr, Y, and Zr
is thought to be synthesized through the s-process in lowto-intermediate mass AGB stars. In the lower metallicity, it
has been suggested that another nucleosynthesis component isneeded to explain observed abundance patterns of extremely
metal-poor stars (e.g., Honda et al. 2006). We discuss this pointby taking the ratio of these elements to Ba in Section 4.1.7.
4.1.6. Heavy Neutron-capture Elements: Ba, La, Nd, Sm, and Eu

Similar to the behavior seen in the [Y/Fe]–[Fe/H] diagram,two of the sample stars, G 18–24 and BD + 04◦ 2466, stand out
in abundance ratios for Ba, La, Nd, and Sm, which may suggestthat surface composition has been modified during their internal
evolution.
All of the three subsamples show a similar [Ba/Fe] trendwith [Fe/H]; [Ba/Fe] is near-solar in the range [Fe/H] > −1.5,
while it decreases toward lower [Fe/H]. The near-solar [Ba/Fe]
ratios for the thick disk stars are in agreement with previous
studies (Prochaska et al. 2000; Reddy et al. 2006). The observed
decreasing [Ba/Fe] with decreasing [Fe/H] observed for thehalo stars is also consistent with the trend reported by Gratton& Sneden (1994) and François et al. (2007).
For the [La/Fe], [Nd/Fe], and [Sm/Fe] ratios, a hint of adifferent trend with [Fe/H] between the thick disk and theinner/outer halo subsamples is modestly seen; the abundance
ratios for the thick disk stars show a flat trend with [Fe/H] atvalues close to the solar ones, while the ratios for the inner/outer
halo stars seem to increase with increasing [Fe/H]. In particular,
at the highest metallicities, the inner/outer halo subsamples
show clear overabundance at 0.42–0.60 dex on average, which
is not seen in the thick disk subsample (see Table 5). In the studyof Mashonkina et al. (2004), who performed Nd abundance
analysis of 60 thin disk, thick disk, and halo stars, using aspectral synthesis techniques, such an abundance difference
between the thick disk and halo stars was not identified. Instead,they reported that a similar overabundance of [Nd/Fe] in therange 0.23–0.45 dex for both thick disk and halo stars exists intheir overlapping metallicity range. We note that some offset inabundances between the two studies exists since Mashonkina
et al. (2004) employed empirical log gf values while this study
adopts the log gf values from Ivans et al. (2006), which wereoriginally taken from the measurement by Den Hartog et al.
(2003). Further precise studies on the heavy neutron-capture
elemental abundances between the thick disk and halo stars
in their overlapping metallicity range with a larger sample aredesirable to convincingly conclude the abundance differencesamong the different Galactic populations.

Figure 9. [Sr/Ba], [Y/Ba], [Zr/Ba], [Nd/Ba], and [Sm/Ba] ratios plottedagainst [Ba/H]. The symbols are the same as in Figure 7. The abundance ratiosof the solar-system s-process (dotted) and r-process (dash-dotted) components
predicted by Arlandini et al. (1999) are indicated by horizontal lines.
(A color version of this figure is available in the online journal.)

previously been reported by François et al. (2007), in which it
was interpreted as a possible indication that an extra production
mechanism for lighter neutron-capture elements over heavier
ones is required. The mechanism likely responsible for theexcess of light neutron-capture elements is frequently called theLight Element Primary Process. Their nature and astrophysicalsites are investigated by nucleosynthesis calculations in corecollapse SNe (Wanajo et al. 2011; Arcones & Montes 2011).
Our results imply that the mechanisms responsible for theSr, Y, and Zr enrichment were commonly efficient among theprogenitor of the thick disk, inner halo, and outer halo stars with
[Ba/H] < −2.0.
In [Ba/H] above ∼−1.5, [Nd/Ba] and [Sm/Ba] ratios fora majority of the thick disk stars and some inner/outer halostars are below the solar-system r-process values. This result is
consistent with the previous study of Mashonkina et al. (2004).
A possible indication is that the s-process may have contributed
to the chemical enrichment in progenitors of both the thick disk
and inner/outer halo components. In the next subsection, we
will discuss the r- and s-process contribution to the observed
abundances of neutron-capture elements in our sample based on
the Eu abundances.

4.1.7. [X/Ba]–[Ba/H] Diagrams for Neutron-capture Elements

Abundance ratios among neutron-capture elements arewidely used as a diagnostic for possible nucleosynthesis sitesof these elements (e.g., François et al. 2007). Figure 9 shows[Sr/Ba], [Y/Ba], [Zr/Ba], [Nd/Ba], and [Sm/Ba] ratios plottedagainst [Ba/H] for the thick disk, inner halo, and outer halosubsamples. The values for the solar-system s- and r-processcomponents from Arlandini et al. (1999) are indicated by thedotted and dash-dotted lines, respectively.
In the range [Ba/H] < −2.0, abundance ratios of thelighter neutron-capture elements, Sr, Y, and Zr, for most of thesample stars show higher values than those for the solar-systemr-process component and an increasing trend toward lower[Ba/H]. Both the inner halo and the outer halo stars, as well as
thick disk stars, seem to follow the above trend. Such a trend has

4.1.8. Europium

The bottom right panels of Figure 8 show the [Eu/Fe]
abundance ratios plotted against [Fe/H] for the thick disk, inner
halo, and outer halo subsamples. The most remarkable featurein these plots is that the [Eu/Fe] ratios are more enhanced in the
inner/outer halo subsamples than in the thick disk subsample
by ∼0.2–0.3 dex in the metallicity range of [Fe/H] > −1.5.
In Table 5, the mean values for the three subsamples taking
into account only dwarfs (μd ) or giants (μg ) are also indicated.
The μd for [Eu/Fe] in the inner and outer halo subsamples is
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Figure 10. Same as Figure 7, but for [Eu/Fe], [Eu/Ba], and [Eu/La]. The
abundance ratios of the solar-system s-process (dotted) and r-process (dashdotted) components predicted by Arlandini et al. (1999) are indicated byhorizontal lines.
(A color version of this figure is available in the online journal.)

Figure 11. [Eu/Mg] ratios plotted against [Mg/H]. The symbols are the sameas in Figure 7.
(A color version of this figure is available in the online journal.)

again larger than that in the thick disk subsample at metallicities[Fe/H] > −1.5, while differences in μg among the threesubsamples are not clear. Although further confirmation witha larger sample is desirable for a definite conclusion, thedifference in μd may indicate that the abundance difference
is not totally caused by systematic errors due to the inclusion ofboth dwarfs and giant stars in our sample.
To examine whether the enhanced [Eu/Fe] ratios observedin the inner and outer halo stars are caused by an excess of ther- or s-process contribution to the production of Eu, the middle
and bottom panels of Figure 10 show the [Eu/Ba] and [Eu/La]
ratios in the sample stars as well as the values for the pure r- ands-process components in the solar-system material predicted by
Arlandini et al. (1999). The [Eu/Ba] ratios in many of the inner
halo and outer halo subsamples show a flat trend with [Fe/H]
up to [Fe/H] ∼ −1.5 and do not significantly deviate from thevalue for the solar-system r-process component ([Eu/Ba] ∼0.7).
In the same metallicity range, the thick disk subsample, on theother hand, predominantly shows the [Eu/Ba] ratios below theflat inner/outer halo trend.
Similar behavior can be seen in the [Eu/La] versus [Fe/H]
plot. Again, in the metallicity range of [Fe/H] > −1.5, thethick disk stars tend to show lower [Eu/La] ratios than the
typical inner/outer halo stars. Some of the inner/outer halostars including the known s-rich stars show much lower [Eu/Ba]
and [Eu/La] ratios close to the values for the solar-systems-process component.
These results suggest that Eu in our inner/outer halo sample
stars is predominantly an r-process origin. In contrast, theobserved lower [Eu/Ba] and [Eu/La] ratios in the thick disk stars
may indicate the contribution of the s-process to the chemicalenrichment during the formation of the thick disk. A similar
conclusion has been reached for the thick disk stars by previous
studies (Mashonkina & Gehren 2000; Bensby et al. 2005).

If Eu is an r-process origin and if the r-process elements
are predominantly ejected in Type II SNe as α elements, it ispuzzling that the trend found in the [Eu/Fe] ratios for these stars
does not follow that seen in the α elements. As can be seen in
the top and middle panels of Figure 11, while the [Mg/Fe] ratios
in many of the inner/outer halo stars follow a decreasing trendwith increasing [Fe/H] (Paper I), the [Eu/Fe] ratios for thesestars do not. This result is in contrast to the expectation that the[Mg/Fe] and [Eu/Fe] ratios would show a similar trend with
[Fe/H] if Mg and Eu are produced in the same astrophysical site,
presumably Type II SNe of massive stars. Rather, the stars withlow [Mg/Fe] ratios (gray circled symbols) in our sample tend
to show higher [Eu/Fe] ratios.
The difference in the behavior of the Eu and Mg abundances
can be more clearly seen in the bottom panel of Figure 11, where
the [Eu/Mg] ratios are plotted against [Mg/H]. In this plot, thethick disk stars show almost a flat [Eu/Mg] trend for all of the[Mg/H] range at about the solar value, while the inner/outer
halo stars exhibit larger values of [Eu/Mg], mildly increasing
with increasing [Mg/H]. The similar [Eu/Mg] enhancement
has been reported previously by Letarte et al. (2010) for stars in
Fornax dwarf spheroidal galaxy. Unlike the present sample, theFornax stars were found to be rich in s-process elements likeBa, and thus the enhanced [Eu/Mg] ratios have been interpreted
as due to the excess of the s-process in this galaxy (Letarte et al.
2010). A different interpretation is likely needed to explain thehigh [Eu/Mg] for the present sample since the signature ofsignificant s-process enrichment like that seen in Fornax is notobserved in these stars as mentioned above.
The implications of these results are twofold: (1) on theastrophysical site of the r-process and (2) on the different
chemical enrichment histories in the progenitors of the thickdisk and the inner/outer stellar halos. For point (1), the results
favor the scenario that the astrophysical site which produces
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It can be seen that in the metallicity below [Fe/H] ∼ −2.0,
stars with any of the extreme orbital parameters will overlap
with the typical inner halo distribution. On the other hand, at
higher metallicities ([Fe/H]  −2), deviation from the typicalinner halo distribution for some elements is identified. Both stars
with extreme rotational velocities and large apocentric distanceshow signatures of lower [Mg/Fe], [Na/Fe], [Ni/Fe], and[Zn/Fe] ratios than the other stars in [Fe/H] > −2.0. The stars
with orbital eccentricity close to unity seem to be more similar
to the normal halo stars.
These results imply a possible correlation of chemical abundances with some of the orbital parameters in [Fe/H] −2.0, while no correlation is expected in the lower [Fe/H].
The abundance–kinematics correlation in the metallicity range−1.5 < [Fe/H] < −0.5 is examined in the next subsection.

Eu is different from that for Mg. Magnesium is thought tobe largely synthesized in massive stars during the hydrostatic
burning and ejected through Type II SNe, while more than 90%
of Eu in the solar-system material is thought to be synthesizedin the r-process (Arlandini et al. 1999), whose astrophysicalsite is still unknown. A Type II SN of a relatively low-massprogenitor (8–10 M ) is suggested to be the primary site for the
r-process (e.g., Wheeler et al. 1998). These low-mass SNe as
well as higher mass >20 M SNe are considered in the Galactic
chemical evolution model and are shown to be consistent
with observations of metal-poor stars (Ishimaru & Wanajo
1999). Alternatively, neutron-star mergers are also suggested
as a possible r-process site (Freiburghaus et al. 1999). Eachof these scenarios suffers from large theoretical uncertainties
and thus the r-process site remains controversial (e.g., Argast
et al. 2004). The observed signature of the different trends in
the [Eu/Fe] and the [Mg/Fe] ratios appear to be consistent
with the scenario that Type II SNe of different progenitor
mass ranges are responsible for the production of dominant Eu(8–10 M ) and Mg (e.g., >10 M ), which was also suggested
from chemical evolution models (e.g., Tsujimoto & Shigeyama
1998). However, the high-[Eu/Mg] ratios at [Fe/H] > −1.5
seen in the inner/outer halo stars may require that the Eu
production have a much longer timescale than Mg. Whether
an alternative scenario such as neutron-star mergers can explain
the abundance results must be tested by chemical evolution
models taking into account the star formation history in possible
progenitor systems of the thick disk and stellar halo.
For point (2), on the chemical enrichment histories in theprogenitors of the thick disk and stellar halo, Paper I showedthat the [Mg/Fe] ratios in the inner/outer halo stars show adecreasing trend with [Fe/H] and lower than in the thick diskstars. Low [Mg/Fe] ratios in stars with [Fe/H] > −1.5 are
frequently interpreted as the result of a larger contribution of
Fe from Type Ia SNe in these stars, presumably due to thelower star formation rate in their progenitor systems. However,rather enhanced [Eu/Fe] ratios found in the low-[Mg/Fe] stars
in our sample may indicate difficulty in interpreting theirabundances by this simple picture alone. This is because thelarger contribution of Fe would also reduce the [Eu/Fe] ratios
in these low-[Mg/Fe] stars than in the thick disk stars, whichis in contrast to the observed trend. Chemical evolution models
that allow variation of not only star formation rates, but also
other factors such as an IMF may be necessary to explain theobserved behavior of [Eu/Fe] and [Mg/Fe] ratios consistently
among the thick disk and halo stars.

4.2.2. Abundance Ratios versus Orbital Parameters

In the previous subsections, we show the differences inelemental abundances in the range [Fe/H] > −2.0 between
stars with extreme kinematics and the stars having typical inner
halo kinematics. In order to examine whether the observed
differences came from any correlation between [X/Fe] withthe orbital parameters among the inner and outer halo stars,
we calculate the linear correlation coefficient for the [X/Fe]with Vφ , Zmax , Rapo , and orbital eccentricity e and probability at
which a null hypothesis of no correlation can be rejected. Someof the [X/Fe] ratios appear to be weekly correlated with [Fe/H].
In order to reduce the effect of the [X/Fe] versus [Fe/H]
correlation in examining the [X/Fe] versus orbital parametercorrelation, we limit the sample stars to those with metallicity−1.5 < [Fe/H] < −0.5. Elements for which more than 20 stars
are available in the above metallicity range are considered. As
a result, we found that for Na, Sc, Ni, and Zn, the possible
correlation is suggested for one or more of the above orbitalparameters with the probability for the null hypotheses of lessthan 5%.
Figure 13 plots (top to bottom) [Na/Fe], [Sc/Fe], [Ni/Fe],and [Zn/Fe] as a function of (left to right) Vφ , log Zmax ,log Rapo , and e for the sample stars with −1.5 < [Fe/H]
 −0.5. As in Figure 7, the crosses, filled circles, andfilled triangles indicate the thick disk stars, the inner
halo stars, and the outer halo stars, respectively. Theirintermediate populations are shown in the open symbols(thick disk/inner halo: open circles; inner halo/outer halo: open
triangles).
In the [X/Fe] versus Vφ plot, the stars on retrograde orbit(Vφ < 0.0) tend to show lower abundance ratios compared to thestars with prograde orbits including thick disk stars (cross) andthick disk/inner halo intermediate stars (open circles). One outer
halo star with an extreme prograde rotation (Vφ ∼ 280 km s−1 )clearly shows a different abundance pattern compared to thethick disk stars.
In the [X/Fe] versus log Zmax plot, the abundance scatter
appears to be larger in Zmax > 1 kpc. The Zmax ∼ 1 kpc
corresponds to the transition between the thick disk dominant
to the inner halo dominant region according to Carollo et al.
(2010). At larger Zmax , the trend is not very clear for all of theelemental abundances.
In the [X/Fe] versus log Rapo plot, stars with Rapo > 30 kpc
clearly show lower [Na/Fe] and [Ni/Fe] ratios compared to theother halo stars with smaller Rapo .
Interestingly, the correlation between the [X/Fe] and e can beseen for all of the four elements. The calculated probability for4.2. Abundance Kinematics Correlation
4.2.1. Abundance Patterns of Stars with Extreme Kinematics

Figure 12 highlights the abundances of stars with extreme
orbital parameters: i.e., stars having extreme rotational velocities(Vφ < −150 or >250 km s−1 : stars), high maximum distancefrom the Galactic plane (Zmax > 20 kpc: crosses), large
apocentric distance (Rapo > 30 kpc: diamonds), and eccentricity
close to unity (e > 0.95: asterisks). Stars that meet more than
one of the four categories are shown with overlapping symbols.
Mean abundances of the typical inner halo stars with PIH > 0.95
in each metallicity bin of 0.5 dex are shown by the solid lines,
where the dotted lines show a range in mean ± dispersion if
there are 2 stars in each [Fe/H] interval. The extreme cut of
PIH > 0.95 was employed to exclude possible contamination ofstars with thick disk-like or outer-halo-like kinematics.
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Figure 12. [X/Fe]–[Fe/H] diagram for stars with extreme kinematics: stars with extreme rotational velocities (Vφ < −150 or >250 km s−1 ; stars), high maximum
distance from the Galactic plain (Zmax > 20 kpc; crosses), large apocentric distance (Rapo > 30 kpc; diamonds), and orbital eccentricity close to unity (e > 0.95;
asterisks). The other sample stars are shown in gray symbols. The solid line in each panel connects mean abundance of the inner halo stars with PIH > 0.95 in each of
the 0.5 dex [Fe/H] intervals. The dotted lines indicate the range within 1σ scatter about the means.
(A color version of this figure is available in the online journal.)

the [X/Fe]–[Fe/H] diagrams for each of the three subsamples
are investigated as a clue to the formation mechanisms of theseold Galactic components. Our results show that the abundances
are largely similar among the thick disk, inner halo, andouter halo subsamples in the metallicity range of [Fe/H] <−2. In contrast, the abundance differences for some elements
among the three subsamples are identified at higher metallicities([Fe/H]  −1.5). Our main results are summarized as follows.
1. The inner halo and outer halo stars show lower [Na/Fe],[Ni/Fe], [Cu/Fe], and [Zn/Fe] ratios than the thick diskstars in the metallicity range of [Fe/H] > −1.5. In
particular, the sample stars with relatively low [Mg/Fe]
show low [Zn/Fe] ratios compared to the other halo stars,
which is in line with the results of Nissen & Schuster (2011).
2. All of the three subsamples show an increasing [Mn/Fe]
trend with [Fe/H], which may suggest that Mn fromType Ia SNe has contributed to enriching the progenitorsof the thick disk, inner and outer stellar halos.

the null hypothesis of no correlation is less than 2% for theseelements. As mentioned in the previous subsection, the highest
e stars consist of both the inner and outer halo stars in our
sample. This might partly reflect our selection of high-velocity
stars in the solar neighborhood, which preferentially select higheccentricity stars. The tendency for these stars to have different
chemical properties may result from our selection bias toward
stars accreted from other systems, although such population
may not be dominant in the inner halo region as suggested by
Carollo et al. (2010)
5. SUMMARY AND DISCUSSION
We estimate elemental abundances of sodium, iron-peak,and neutron-capture elements of 97 dwarf and giant stars
kinematically belonging to the MW thick disk, inner halo,
and outer halo components in the metallicity range of −3.3 <[Fe/H] < −0.5. Using these samples, characteristic trends in
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Figure 13. From top to bottom: [Na/Fe], [Sc/Fe], [Ni/Fe], and [Zn/Fe] plotted against orbital parameters. From left to right: Vφ , log Zmax , log Rapo , and orbital
eccentricity in the metallicity range −1.5 < [Fe/H]  −0.5. The mean values and their 1σ scatter in each bin of the orbital parameters are indicated with solid anddotted lines, respectively. The symbols are the same as in Figure 7.
(A color version of this figure is available in the online journal.)

3. The [Eu/Fe] ratios in the inner/outer halo stars are higher
than in the thick disk stars in the range [Fe/H] > −1.5.
This behavior is in contrast to that seen in [Mg/Fe] ratios,
for which many of the inner/outer halo stars show lower
values than the thick disk stars. The different behavior of[Eu/Fe] and [Mg/Fe] ratios for the three subsamples may
imply that the production site of Eu is different from thatof Mg, although more exact interpretations require detailed
chemical evolution modeling in possible progenitors of the
thick disk and stellar halos as well as determination of yields
for these elements.
4. The [Eu/Ba] and [Eu/La] ratios for the thick disk stars and ahandful of the inner halo stars are below the values expected
for the solar-system r-process component (Arlandini et al.
1999), which may indicate the contribution of the s-processfrom low-to-intermediate mass AGB stars to the chemicalenrichment in the progenitors of these stars. On the other
hand, many of the inner and outer halo stars follow theflat trend up to [Fe/H] ∼ −1.5, likely suggesting that the
r-process dominates in synthesizing heavy neutron-capture
elements in their progenitor systems.
5. The inner/outer halo stars with extreme retrograde rotation,
large Zmax , and large Rapo tend to show lower [Na/Fe],[Ni/Fe], and [Zn/Fe] ratios than those with normal inner
halo kinematics in [Fe/H] > −2. In the lower metallicities,
the abundances are almost indistinguishable among thekinematically different populations.

6. In our inner and outer halo subsamples with metallicities−1.5 < [Fe/H]  −0.5, signatures of the correlation
between [Na/Fe], [Sc/Fe], [Ni/Fe], and [Zn/Fe] ratios andthe orbital eccentricity are identified.
Based on the above results, we discuss the possible formation
mechanisms of the thick disk, inner and outer stellar halos,
with particular focus on what are the likely progenitors of thesecomponents.
5.1. The Thick Disk
Several scenarios from theoretical models for the thick disk
formation are proposed; (1) monolithic dissipative collapse ofa disk component (e.g., Burkert et al. 1992), (2) the mergers ofsatellites that are tidally disrupted to populate the thick disk (e.g.,
Abadi et al. 2003), (3) heating a pre-existing thin disk throughnumerous mergers of dark matter subhalos or dwarf galaxies(e.g., Quinn et al. 1993; Velázquez & White 1999; Hayashi &Chiba 2006), (4) multiple dissipative mergers of building blocks
that trigger rapid star formation (e.g., Brook et al. 2004, 2005),
and (5) early evolution of a gas-rich clumpy young disk (e.g.,
Bournaud et al. 2009). It has also been proposed that (6) the
secular radial migration of the thin disk may be responsible
for the thick disk formation (e.g., Haywood 2008; Schönrich &Binney 2009). Which of the above mechanisms is dominant andhow the formation of the thick disk relates to the formation of21

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Ishigaki, Aoki, & Chiba

the other Galactic components, namely, the bulge, thin disk, andstellar halo, are still controversial.
In the present study and in Paper I, we found the remarkable
abundance difference for several elements between the thickdisk and the inner/outer halo stars in their overlapping metallicity range (−1.5  [Fe/H]  −0.5). The main implications fromthis result on the formation mechanisms of the thick disk component are summarized as follows; first, the high-[Na/Fe] ratios,
as well as high-[Mg/Fe] ratios reported in Paper I, comparedto many of the inner/outer halo stars indicate that chemical
evolution of the thick disk was driven primarily through nucleosynthesis products from Type II SNe. This may suggest that the
thick disk formation timescale was sufficiently short and/or the
IMF in the progenitor of the thick disk was biased toward highmass stars so that nucleosynthesis products of low-mass stars(i.e., Type Ia SNe) played a minor role. Second, the increase
in [Mn/Fe] with [Fe/H] for the thick disk stars indicates thatsome contribution of Type Ia SNe may present, since Mn is predominantly synthesized in the explosive burning of Type Ia SNein the metallicity typical of the thick disk stars (Kobayashi &Nakazato 2011). Third, the [Eu/Ba], [Eu/La], [Sm/Ba], and[Nd/Ba] ratios for the thick disk are below the values expected
for the solar-system r-process components, which suggests thatthe s-process has contributed to the chemical enrichment of
the thick disk. Fourth, the lower [Eu/Fe] and [Eu/Mg] ratios
in the thick disk stars than in the inner/outer halo stars further support distinct chemical enrichment histories among thesecomponents, while its interpretation depends on a currently unknown astrophysical site for Eu production. Finally, for most of
the elements, scatter in the abundance ratios for the thick diskstars is comparable to or smaller than the observational errors,
which is in contrast to that for the inner/outer halo stars with
[Fe/H] > −1.5. The lack of scatter may indicate that the gas
from which the thick disk has formed may have been relatively
well mixed.
The short timescale for the formation of the thick disk might
be expected in scenarios (1), (4), and (5), where sufficient cold
gas is supplied in the form of smooth gas accretion (Bournaud
et al. 2009) or multiple gaseous mergers (Brook et al. 2004)
at high redshift. Scenarios (3) and (6) could also be possible
if a pre-existing thin disk formed under the sufficiently highstar formation rate so that the whole disk stars rapidly became
metallicity as high as the thick disk stars. In each of the above
cases, the nucleosynthesis products of Type Ia SNe and s-process
should have been rapidly mixed in the ISM. Consistent withprevious studies, the abundances of the thick disk stars in our
sample do not resemble those of the known dwarf satellite
galaxies around the MW, like Fornax, whose abundance ischaracterized by much lower [α/Fe] and [Na/Fe] and higher[Ba/Fe] at similar [Fe/H] (Letarte et al. 2010). Therefore, it
is unlikely that the thick disk has been totally built throughan assembly of dSphs having similar properties as surviving
Galactic dSphs (e.g., scenario (2)).
The four thick disk stars with [Fe/H] < −0.8 in our sample,
all of which have kinematic properties similar to the proposed
MWTD component, have particular implication on the earliest
evolution of the MW thick disk. As shown in the previous
sections and in Paper I, the candidate MWTD stars show
distinct abundance ratios for several elements from those of
the inner/outer halo stars at metallicities [Fe/H]  −1.5. Even
though the abundance ratios in the candidate MWTD become
almost indistinguishable from the inner/outer halos at lower
metallicities, their trends with [Fe/H] seem to smoothly follow

those seen in the canonical thick disk stars (i.e., the thick diskstars with −0.8 < [Fe/H] < −0.5) with relatively small scatter.
This result implies that these candidate MWTD stars may haveindeed formed within the progenitors of the thick disk, which
are distinct from those of the stellar halo. In other words, ourresult favors the interpretation that they are remnants of theancient MW disk system rather than the interloper of the halo
population that happened to acquire disk-like kinematics. A
possible implication is that whatever the formation mechanism
is, the progenitor of the thick disk may have been a memberof a metal-poor disk system, which later experienced relatively
rapid and homogeneous chemical evolution. Further studies on
the detailed chemical abundances versus kinematics as well asages of individual metal-poor stars are necessary to get deeper
insight into the formation of the thick disk and the origin of its
MWTD component.
5.2. The Stellar Halos
One of the key questions on the global formation andevolution of the MW stellar halos is whether the halos havebeen assembled from smaller stellar systems similar to thedSph galaxies currently orbiting the MW. In the following
subsections, we first address this question by examining whether
the observed abundances in our inner/outer halo subsamples
resemble those reported for stars in classical as well as ultra-faintdwarf galaxies. For this comparison, we separately consider two
metallicity ranges ([Fe/H] > −2 and <−2) since observed
abundance trends for the inner/outer halo subsamples varydepending on their metallicity. Then, we discuss whether theseabundances in the MW thick disk, inner and outer stellar halos
are compatible with proposed MW formation scenarios basedon other observations or numerical simulations.
5.2.1. Comparisons with MW dSph Galaxies[Fe/H] < −2. In metallicities below −2.0, the abundancesof several elements in both the inner and outer halo subsamples
are found to be in general agreement with known classicaland ultra-faint dwarf galaxies regardless of their kinematics.
In this metallicity range, the inner/outer halo stars generally
show enhanced [Mg/Fe] and [Si/Fe] ratios, near-solar ratios
for [Ni/Fe] or [Sc/Fe], and the low [Mn/Fe] ratios. Similarabundance ratios were reported for the lowest metallicity stars([Fe/H]  −2.5) in classical dwarf galaxies like Draco (Cohen& Huang 2009), Ursa minor (Cohen & Huang 2010), Carina(Venn et al. 2012), and Sculptor (Starkenburg et al. 2013). The
abundances of light elements (Z < 30) in newly discovered
ultra-faint dwarf galaxies also show abundances similar to thehalo stars (Frebel et al. 2010; Simon et al. 2010; Gilmore et al.
2013) in the whole metallicity range spanned by these stars([Fe/H]  −2.0).
Despite these similarities, our sample of both inner and outer
halo stars with [Fe/H] < −2.0 also confirms some of the known
discrepancies in abundances between the dSphs and MW halostars. First, as discussed in Paper I, our sample of inner and outer
halo stars both shows enhanced [Mg/Fe] ratios, while those in
extremely metal-poor stars in Sextans dSph are systematicallylower than the MW halo stars at similar metallicities (Aoki et al.
2009). Second, abundance ratios of neutron-capture elements
like [Sr/Fe] or [Ba/Fe] tend to be low in both classical (Aokiet al. 2009; Venn et al. 2012) and ultra-faint dwarf galaxies(Koch et al. 2008; Frebel et al. 2010; Simon et al. 2010), with
22

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Ishigaki, Aoki, & Chiba

defined outer halo stars with [Fe/H] > −2.0 could be a relatively minor population (see Section 5.2.2). The implicationfrom the abundance results in the present study and in Paper I isthat the progenitors of these relatively metal-rich outer halo stars
may have experienced a certain interval of chemical evolutionthat has allowed the system to be enriched with nucleosynthesis
products of low-to-intermediate mass stars. It is clear that thecurrently known ultra-faint dwarf galaxies cannot be a dominant
progenitor in metallicities [Fe/H] > −2, since they are moremetal-poor and tend to show different chemical abundance patterns. Indeed, the known ultra-faint dwarf galaxies were reported
to have much simpler stellar population which have ended theirmajor star formation events at much earlier times than brighter
dwarf galaxies (Okamoto et al. 2012). Although the surviving
dSphs like Fornax or Sagittarius may not be a direct progenitor of these stars, the similarity in light elements may suggestthat gas-rich systems that have longer star formation timescales
have contributed to the relatively metal-rich part of the outer
stellar halo.

an exception of a possible binary star (Honda et al. 2011),
while the field halo stars including our inner/outer halo samples
exhibit both high and low abundances of these neutron-capture
elements. In conclusion, the abundance differences as indicatedabove remain even if the present sample of the outer halo stars
with extreme kinematics is taken into account. The suggested
abundance differences in neutron-capture elements may point
to some differences in chemical evolution between progenitorsof the MW field halo and some of the dSphs. In order toget deeper insights into this discrepancy, chemical evolution
modeling which takes into account star formation and mixingof elements within a small system as well as investigation of
nucleosynthesis yields of neutron-capture elements at this low
metallicity would be necessary.
To summarize, in the metallicity below [Fe/H] ∼ −2.0,
abundances of α and Fe peak elements in both of our innerand outer halo subsamples largely overlap with those of theextremely metal-poor stars in some of the classical and ultrafaint dwarf galaxies. The similarity supports the hypothesisthat metal-poor stars in both of the inner and outer halos inthis low-metallicity range are at least in part accreted fromsystems similar to many of the currently known classical andultra-faint dwarf galaxies. We note that, as we have mentioned
in Section 2.2, recent surveys suggest that the metallicity
distribution function is different between the inner and outer
halos, and thus, the inner halo stars with [Fe/H] ∼ −2.0 are
relatively minor population. Therefore, the accretions of suchmetal-poor dSphs-like systems are likely more dominant in theouter halo than in the inner halo.
[Fe/H] > −2. In this metallicity range, abundances of several
elements in most of the outer halo stars, especially those withextreme kinematics, overlap with those in the MW dSph stars,
while the stars with the typical inner halo kinematics typicallyshow a larger discrepancy in the abundances from the dSphs
as recognized in previous studies (Venn et al. 2004; Tolstoy
et al. 2009). As an example, [Na/Fe] ratios in stars in Fornax
and Sagittarius dSphs with −1.5 < [Fe/H] < −0.5 have beenreported to be subsolar in a range −1 < [Na/Fe] < −0.2
(Letarte et al. 2010; Sbordone et al. 2007), while the present
sample of the outer halo stars also shows subsolar [Na/Fe]
ratios in contrast to the thick disk and some of the inner halostars mostly showing the supersolar ratios. A similar trend isfound in the [Ni/Fe] ratios, for which the outer halo stars largely
show subsolar values similar to stars in Fornax and Sagittarius
dSphs (Letarte et al. 2010; Sbordone et al. 2007), while the thickdisk and some of the inner halo stars generally show supersolar
values.
On the other hand, differences in Ba abundances in thesedSph stars (Sbordone et al. 2007; Letarte et al. 2010) andthe inner/outer halo stars seem to be substantial regardless ofkinematics. Ba in moderately metal-poor stars is thought to bemainly produced in low-to-intermediate mass AGB stars. The
difference between the MW inner/outer halo stars and the dSphs
indicates that chemical enrichment by low-to-intermediate mass
stars may be different in the progenitors of the inner/outer halofrom the surviving dSphs.
Then, the intriguing question is: what is the origin of thekinematically classified outer halo stars with [Fe/H] > −2.0
which show chemical abundance ratios different from the normal inner halo stars? Again, we note that recent determinationof the metallicity distribution function (e.g., An et al. 2013)
suggests that the outer halo population is dominated by moremetal-poor stars ([Fe/H] ∼ −2.0), and thus our kinematically

5.2.2. Comparisons with Other Observations
and Numerical Simulations

In order to accommodate the observed phase space and chemical abundance distributions, the formation of the stellar halo isbelieved to have involved multiple processes. These processes
are frequently classified into two mechanisms, namely, a dissipational collapse of a proto-Galactic gas cloud within a very
short timescale, which was introduced by a landmark study of
Eggen et al. (1962), and dissipationless mergers of smaller stellar systems with much longer timescales, which was proposed
by Searle & Zinn (1978). The existing observations of halo stars
seem to favor the scenario that the inner part of the stellar halo isformed via dissipative collapse of gaseous proto-galactic fragments while the outer part has built up through dissipationlessmerging of small stellar systems (e.g., Chiba & Beers 2000;
Carollo et al. 2007, 2010). Recently, An et al. (2013) studied metallicity of the stellar halo at heliocentric distances in therange 5–8 kpc based on the SDSS ugriz photometry. They foundthat the metallicity distribution function can be modeled withtwo Gaussian components with peaks at [Fe/H] ∼ −1.7 and∼−2.3, which favors the above scenario that at least two mechanisms are responsible for building up the present-day stellar
halo. It was also shown that such a hybrid scenario can naturally occur in the context of galaxy formation theory under thecurrent standard cosmology (e.g., Bekki & Chiba 2001; Zolotov
et al. 2010; Font et al. 2011; McCarthy et al. 2012; Tissera et al.
2012, 2013).
Our results show that the inner halo at relatively low metallicity ([Fe/H] < −2) has similar abundance ratios as the outer
halo and is broadly consistent with previous abundance results
for the MW dSphs. This may suggest that, in [Fe/H] < −2, inner
halo stars were, at least in part, assembled from progenitor systems similar to these dSphs. The presence of stars accreted fromdSphs is also predicted from cosmological hydrodynamical simulation of McCarthy et al. (2012), which suggests that, although
the inner halo is expected to be dominated by stars formed insitu within the proto-disk and later dynamically heated up to thehalo region, a small fraction of accreted populations is indeed
expected. In the higher metallicity range, the abundance ratios
in the inner halo subsample for some elements were clearly
different from those in the thick disk subsample. In particular
generally larger dispersion in abundance ratios in the inner halostars compared to those in the thick disk stars argues against
23

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba

the hypothesis that the inner halo component is entirely built up
with a single dissipational collapse and star formation within awell-mixed gas. Rather, the results favor the scenario that multiple gas-rich systems that have experienced various levels ofchemical evolution prior to the merging have contributed to thepresent-day inner halo.
The cosmological simulation of Zolotov et al. (2010) suggeststhat a certain fraction of inner halo stars are expected to form
within the central region of a galaxy later heated up to halolike orbits via mergers and such stars can be distinguished fromaccreted stars with their chemical abundances. Stars that arelikely formed in such a process can indeed frequently be foundin our inner halo subsample with [Fe/H] > −2, characterized
as high [Mg/Fe], [Si/Fe], [Na/Fe], and [Zn/Fe] ratios. The
simulation of McCarthy et al. (2012) suggests that the large
fraction of the inner halo is populated by the stars formed in
the disk system later puffed up by the dynamical heating via
merging dark matter subhalos. The candidate of such population
can be found in our sample stars having kinematics intermediate
between the thick disk and the inner halo subsamples. Thesestars are found to have similar abundance as the thick diskstars inferring that a certain fraction of the halo is populated
by stars originally belonging to the MW disk and later heatedup while partly conserving their initial kinematics and chemical
abundances.
Our outer halo subsample shows broadly similar abundance
ratios of light elements (Z < 30) with many of the MW
dSphs in [Fe/H] < −2.0, which may support the hypothesis that the outer stellar halo, where Carollo et al. (2007,2010) found it to be dominant in [Fe/H] < −2, was assembled from stellar systems similar to the present-day lowluminosity MW dSphs. On the other hand, our outer halo subsample with higher metallicity ([Fe/H] > −2) shows different
abundance ratios for some elements compared to the inner halostars. Our sample of more metal-rich outer halo stars is generally similar to the low-α stars reported by Nissen & Schuster (2010, 2011) in terms of abundances (e.g., low [Na/Fe]
and [Zn/Fe]) and kinematics (e.g., extreme retrograde orbit),although distinction from higher-α halo stars at a given [Fe/H]
is not as clear as that observed in their works. These results
may suggest that systems that have experienced extended period of chemical evolution including Type Ia SNe before theiraccretion may have contributed to relatively metal-rich outer
halo stars in the solar neighborhood that have distinct abundances compared to the bulk of inner halo stars. The contribution of a certain fraction of the accreted stars is also suggested
from medium resolution spectroscopic samples that include
more distant halo stars (Schlaufman et al. 2012; Sheffield et al.
2012).
The present study of chemical abundances in the kinematically selected sample stars provides important insights into theprogenitor systems that gave rise to the present stellar halo atleast in the solar neighborhood. However, the number of sample stars is too small and incomplete to quantify the fractional
contribution of accretions to build up the stellar halo in the solar
neighborhood. Upcoming surveys such as Gaia and its spectroscopic followup will provide valuable insights into the origins
of halo stars by measuring accurate phase-space informationand chemical abundances of a large number of stars withinseveral kpc from the Sun. Besides, kinematics and chemical
abundance analysis of in situ outer halo stars at Galactocentric
distance of several tens of kpc are important to constrain thenature of the outer stellar halo.

We are grateful to the referee for a careful reading of thismanuscript and for a number of constructive comments, whichhave significantly improved our paper. We thank A. Tajitsu,
T.-S. Pyo, and the staff members of the Subaru telescope for theirhelpful support and assistance with our HDS observation. M.N.I.
is grateful to S. Wanajo and T. Tsujimoto for useful discussions
and comments. This work is supported in part by Grant-in-Aid
for Scientific Research (23740162,23224004) of the Ministry of
Education, Culture, Sports, Science, and Technology of Japan.
Facilities: Subaru (HDS).
REFERENCES
Abadi, M. G., Navarro, J. F., Steinmetz, M., & Eke, V. R. 2003, ApJ, 597, 21
Alves-Brito, A., Meléndez, J., Asplund, M., Ramı́rez, I., & Yong, D. 2010, A&A,
513, 35
An, D., Beers, T. C., Johnson, J. A., et al. 2013, ApJ, 763, 65
Anders, E., & Grevesse, N. 1989, GeCoA, 53, 197
Andrievsky, S. M., Spite, M., Korotin, S. A., et al. 2007, A&A, 464, 1081
Aoki, W., Barklem, P. S., Beers, T. C., et al. 2009, ApJ, 698, 1803
Arcones, A., & Montes, F. 2011, ApJ, 731, 5
Argast, D., Samland, M., Thielemann, F.-K., & Qian, Y.-Z. 2004, A&A,
416, 997
Arlandini, C., Käppeler, F., Wisshak, K., et al. 1999, ApJ, 525, 886
Asplund, M., Grevesse, N., Sauval, A. J., & Scott, P. 2009, ARA&A, 47, 481
Beers, T. C., Carollo, D., Ivezić, Ẑ., et al. 2012, ApJ, 746, 34
Beers, T. C., Chiba, M., Yoshii, Y., et al. 2000, AJ, 119, 2866
Beers, T. C., Drilling, J. S., Rossi, S., et al. 2002, AJ, 124, 931
Bekki, K., & Chiba, M. 2001, ApJ, 558, 666
Bell, E. F., Zucker, D. B., Belokurov, V., et al. 2008, ApJ, 680, 295
Bensby, T., Feltzing, S., Johnson, J. A., et al. 2010, A&A, 512A, 41
Bensby, T., Feltzing, S., & Lundström, I. 2003, A&A, 410, 527
Bensby, T., Feltzing, S., Lundström, I., & Ilyin, I. 2005, A&A, 433, 185
Blackwell-Whitehead, R., & Bergemann, M. 2007, A&A, 472, L43
Bonifacio, P., Spite, M., Cayrel, R., et al. 2009, A&A, 501, 519
Bournaud, F., Elmegreen, B. G., & Martig, M. 2009, ApJL, 707, L1
Bovy, J., Rix, H-W., & Hovv, D. W. 2012, ApJ, 751, 131
Brook, C. B., Gibson, B. K., Martel, H., & Kawata, D. 2005, ApJ, 630, 298
Brook, C. B., Kawata, D., Gibson, B. K., & Freeman, K. C. 2004, ApJ, 612, 894
Burkert, A., Truran, J. W., & Hensler, G. 1992, ApJ, 391, 651
Burris, D. L., Pilachowski, C. A., Armandroff, T. E., et al. 2000, ApJ, 544, 302
Carney, B. W., Latham, D. W., Laird, J. B., & Aguilar, L. A. 1994, AJ,
107, 2240
Carollo, D., Beers, T. C., Chiba, M., et al. 2010, ApJ, 712, 692
Carollo, D., Beers, T. C., Lee, Y. S., et al. 2007, Natur, 450, 1020
Casagrande, L., Ramı́rez, I., Meléndez, J., Bessell, M., & Asplund, M.
2010, A&A, 512, 54
Castelli, F., & Kurucz, R. L. 2003, IAU Symp., Modelling of Stellar Atmosphereed. N. Piskunov (Cambridge: Cambridge Univ. Press)
Cayrel, R., Depagne, E., Spite, M., et al. 2004, A&A, 416, 1117
Cescutti, G., Matteucci, F., Lanfranchi, G. A., & McWilliam, A. 2008, A&A,
491, 401
Cheng, J. Y., Rockosi, C. M., Morrison, H. L., et al. 2012, ApJ, 752, 51
Chiba, M., & Beers, T. C. 2000, AJ, 119, 2843
Cohen, J. G., & Huang, W. 2009, ApJ, 701, 1053
Cohen, J. G., & Huang, W. 2010, ApJ, 719, 931
De Jong, J. T. A., Yanny, B., Rix, H-W., et al. 2010, ApJ, 714, 663
Den Hartog, E. A., Lawler, J. E., Sneden, C., & Cowan, J. J. 2003, ApJS,
148, 543
Eggen, O. J., Lynden-Bell, D., & Sandage, A. R. 1962, ApJ, 136, 748
Font, A. S., McCarthy, I. G., Crain, R. A., et al. 2011, MNRAS, 416, 2802
François, P., Depagne, E., Hill, V., et al. 2007, A&A, 476, 935
Frebel, A., Simon, J. D., Geha, M., & Willman, B. 2010, ApJ, 708, 560
Freeman, K., & Bland-Hawthorn, J. 2002, ARA&A, 40, 487
Freiburghaus, C., Rosswog, S., & Thielemann, F.-K. 1999, ApJL, 525, L121
Fuhr, J. R., & Wiese, W. L. 2005, in NIST Atomic Transition Probability Tables,
Vol. 78, CRC Handbook of Chemistry and Physics, ed. D. R. Lide (Boca
Raton, FL: CRC Press), 10
Fuhrmann, K. 1998, A&A, 338, 16
Fulbright, J. P. 2002, AJ, 123, 404
Gilmore, G., Norris, J. E., Monaco, L., et al. 2013, ApJ, 763, 61
Gilmore, G., & Reid, N. 1983, MNRAS, 202, 1025
Gratton, R. G., & Sneden, C. 1991, A&A, 241, 510
Gratton, R. G., & Sneden, C. 1994, A&A, 287, 927

24

The Astrophysical Journal, 771:67 (25pp), 2013 July 1

Ishigaki, Aoki, & Chiba
Okamoto, S., Arimoto, N., Yamada, Y., & Masao, O. 2012, ApJ, 744, 96
Pickering, J. C. 1996, ApJS, 107, 811
Pignatari, M., Gallino, R., Heil, M., et al. 2010, ApJ, 710, 1557
Prochaska, J. X., Naumov, S. O., Carney, B. W., McWilliam, A., & Wolfe,
A. M. 2000, ApJ, 120, 2513
Quinn, P. J., Hernquist, L., & Fullagar, D. P. 1993, ApJ, 403, 74
Ramı́rez, I., & Meléndez, J. 2005, ApJ, 626, 465
Reddy, B. E., & Lambert, D. L. 2008, MNRAS, 391, 95
Reddy, B. E., Lambert, D. L., & Allende Prieto, C. 2006, MNRAS, 367, 1329
Reddy, B. E., Tomkin, J., Lambert, D. L., & Allende Prieto, C. 2003, MNRAS,
340, 304
Roederer, I. U. 2008, AJ, 137, 272
Roederer, I. U., Sneden, C., Thompson, I. B., Preston, G. W., & Shectman, S.
A. 2010, ApJ, 711, 573
Ruchti, G. R., Fulbright, J. P., Wyse, R. F. G., et al. 2011, ApJ, 737, 9
Ryan, S. G., & Norris, J. E. 1991, AJ, 101, 1835
Sbordone, L., Bonifacio, P., Buonanno, R., et al. 2007, A&A, 465, 815
Schlafly, E. F., & Finkbeiner, D. P. 2011, ApJ, 737, 103
Schlaufman, K. C., Rockosi, C. M., Allende Prieto, C., et al. 2009, ApJ,
703, 2177
Schlaufman, K. C., Rockosi, C. M., Lee, Y. S., et al. 2012, ApJ, 749, 77
Schönrich, R., Asplund, M., & Casagrande, L. 2011, MNRAS, 415, 3807
Schönrich, R., & Binney, J. 2009, MNRAS, 396, 203
Searle, L., & Zinn, R. 1978, ApJ, 225, 357
Sheffield, A. A., Majewski, S. R., Johnston, K., et al. 2012, ApJ, 761, 161
Simmerer, J., Sneden, C., Cowan, J. J., et al. 2004, ApJ, 617, 1091
Simmerer, J., Sneden, C., Ivans, I. I., et al. 2003, AJ, 125, 2018
Simon, J. D., Frebel, A., McWilliam, A., Kirby, E. N., & Thompson, I. B.
2010, ApJ, 716, 446
Sobeck, J. S., Lawler, J. E., & Sneden, C. 2007, ApJ, 667, 1267
Starkenburg, E., Hill, V., Tolstoy, E., et al. 2013, A&A, 549, 88
Stephens, A., & Boesgaard, A. M. 2002, AJ, 123, 1647
Takeda, Y., Hashimoto, O., Taguchi, H., et al. 2005, PASJ, 57, 751
Takeda, Y., Zhao, G., Takada-Hidai, M., et al. 2003, CJAA, 4, 316
Timmes, F. X., Woosley, S. E., & Weaver, T. A. 1995, ApJS, 98, 617
Tissera, P. B., Scannapieco, C., Beers, T., & Carollo, D. 2013, MNRAS, 1371
Tissera, P. B., White, S. D. M., & Scannapieco, C. 2012, MNRAS, 420, 255
Tolstoy, E., Hill, V., & Tosi, M. 2009, ARA&A, 47, 371
Totani, T., Morokuma, T., Oda, T., Doi, M., & Yasuda, N. 2008, PASJ,60, 1327
Tsujimoto, T., & Shigeyama, T. 1998, ApJL, 508, L151
Umeda, H., & Nomoto, K. 2005, ApJ, 619, 427
Unsöld, A. 1955, Physik der Sternatmosphären, (Berlin: Springer)
Velázquez, H., & White, S. D. M. 1999, MNRAS, 304, 254
Venn, K. A., Irwin, M., Shetrone, M. D., et al. 2004, AJ, 128, 1177
Venn, K. A., Shetrone, M. D., Irwin, M., et al. 2012, ApJ, 751, 102
Wanajo, S., Janka, H-T., & Müller, B. 2011, ApJL, 726, L15
Wheeler, J. C., Cowan, J. J., & Hillebrandt, W. 1998, ApJL, 493, L101
Woosley, S. E., & Weaver, T. A. 1995, ApJS, 101, 181
Wyse, R. F. G., & Gilmore, G. 1995, AJ, 110, 2771
Yoshii, Y. 1982, PASJ, 34, 365
Zhang, L., Ishigaki, M., Aoki, W., Zhao, G., & Chiba, M. 2009, ApJ, 706, 1095
Zhao, G., & Magain, P. 1990, A&A, 238, 242
Zolotov, A., Willman, B., Brooks, A., et al. 2010, ApJ, 721, 738

Hayashi, H., & Chiba, M. 2006, PASJ, 58, 835
Haywood, M. 2008, MNRAS, 388, 1175
Helmi, A. 2008, A&ARv, 15, 145
Helmi, A., White, S. D. M., de Zeeuw, T., & Zhao, H. 1999, Natur, 402, 53
Honda, S., Aoki, W., Arimoto, N., & Sadakane, K. 2011, PASJ, 63, 523
Honda, S., Aoki, W., Ishimaru, Y., Wanajo, S., & Ryan, S. G. 2006, ApJ,
643, 1180
Honda, S., Aoki, W., Kajino, T., et al. 2004, ApJ, 607, 474
Ibata, R. A., Gilmore, G., & Irwin, M. J. 1994, Natur, 370, 194
Ishigaki, M., Chiba, M., & Aoki, W. 2010, PASJ, 62, 143
Ishigaki, M., Chiba, M., & Aoki, W. 2012, ApJ, 753, 64
Ishimaru, Y., & Wanajo, S. 1999, ApJL, 511, L33
Ivans, I. I., Simmerer, J., Sneden, C., et al. 2006, ApJ, 645, 613
Ivezić, Ž., Beers, T. C., & Jurić, M. 2012, ARA&A, 50, 251
Iwamoto, K., Brachwitz, F., Nomoto, K., et al. 1999, ApJS, 125, 439
Jorissen, A., Začs, L., Udry, S., Lindgren, H., & Musaev, F. A. 2005, A&A,
441, 1135
Jurić, M., Ivezić, Ž., Brooks, A., et al. 2008, ApJ, 673, 864
Kepley, A. A., Morrison, H. L., Helmi, A., et al. 2007, AJ, 134, 1579
Kobayashi, C., & Nakazato, N. 2011, ApJ, 729, 16
Kobayashi, C., Umeda, H., Nomoto, K., Tominaga, N., & Ohkubo, T. 2006, ApJ,
653, 1145
Koch, A., McWilliam, A., Grebel, E. K., Zucker, D. B., & Belokurov, V.
2008, ApJL, 688, L13
Kurucz, R. L. 1995, in ASP Conf. Ser. 81, Workshop on Laboratory andAstronomical High Resolution Spectra, ed. A. J. Sauval, R. Blomme, &N. Grevesse (San Francisco, CA: ASP), 583
Lai, D. K., Bolte, M., Johnson, J. A., et al. 2008, ApJ, 681, 1524
Lanfranchi, G. A., & Matteucci, F. 2003, MNRAS, 345, 71
Latham, D. W., Stefanik, R. P., Torres, G., et al. 2002, AJ, 124, 1144
Lawler, J. E., Wickliffe, M. E., & Den Hartog, E. A. 2001, ApJ, 563, 1075
Lee, Y. S., Beers, T. C., An, D., et al. 2011, ApJ, 738, 187
Letarte, B., Hill, V., Tolstoy, E., et al. 2010, A&A, 523, 17
Luck, R. E., & Bond, H. E. 1981, ApJ, 244, 919
Majewski, S. R., Skrutskie, M. F., Weinberg, M. D., & Ostheimer, J. C.
2003, ApJ, 599, 1082
Maoz, D., Sharon, K., & Gal-Yam, A. 2010, ApJ, 722, 1879
Mashonkina, L., & Gehren, T. 2000, A&A, 364, 249
Mashonkina, L., Ryabtsev, A., & Frebel, A. 2012, A&A, 540, 98
Mashonkina, L. I., Kamaeva, L. A., Samotoev, V. A., & Sakhibullin, N. A. 2004,
ARep, 48, 185
Matteucci, F., & Greggio, L. 1986, A&A, 154, 279
McCarthy, I. G., Font, A. S., Crain, R. A., et al. 2012, MNRAS, 420, 2245
McWilliam, A. 1998, AJ, 115, 1640
McWilliam, A., Preston, G. W., Sneden, C., & Searle, L. 1995, AJ, 109, 2757
Melendez, J., Asplund, M., Alves-Brito, A., et al. 2008, A&A, 484, 21
Mishenina, T. V., Kovtyukh, V. V., Soubiran, C., Travaglio, C., & Busso, M.
2002, A&A, 396, 189
Newberg, H. J., Yanny, B., & Rockosi, C. 2002, ApJ, 569, 245
Nissen, P. E., Chen, Y. Q., Schuster, W. J., & Zhao, G. 2000, A&A, 2000,
353, 722
Nissen, P. E., & Schuster, W. J. 2010, A&A, 511, 10
Nissen, P. E., & Schuster, W. J. 2011, A&A, 530, 15
Noguchi, K., Aoki, W., Kawanomoto, S., et al. 2002, PASJ, 54, 855

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