c ESO 2014

Astronomy & Astrophysics manuscript no. boo16
January 8, 2014

Chemical compositions of six metal-poor stars in the ultra-faintdwarf spheroidal galaxy Boötes I
M. N. Ishigaki1 , W. Aoki2, 3 , N. Arimoto2, 3, 4 , and S. Okamoto5
1

Kavli Institute for the Physics and Mathematics of the Universe (WPI), Todai Institute for Advanced Study, University of Tokyo,
5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8583, Japan e-mail: miho.ishigaki@ipmu.jp
2
National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan e-mail: aoki.wako@nao.ac.jp,
arimoto.n@nao.ac.jp
3
Department of Astronomical Science, The Graduate University of Advanced Studies, Mitaka, Tokyo 181-8588, Japan
4
Subaru Telescope, National Astronomical Observatory of Japan, 650 North A’ohoku Place, Hilo, HI 96720, USA
5
Kavli Institute for Astronomy and Astrophysics, Peking University, Beijing 100871, China e-mail: okamoto@pku.edu.cn

ABSTRACT

Context. Ultra-faint dwarf galaxies recently discovered around the Milky Way (MW) contain extremely metal-poor stars, and might
represent the building blocks of low-metallicity components of the MW. Among them, the Boötes I dwarf spheroidal galaxy is ofparticular interest because of its exclusively old stellar population. Detailed chemical compositions of individual stars in this galaxy
are a key to understanding formation and chemical evolution in the oldest galaxies in the Universe and their roles in building up theMW halo.
Aims. Previous studies of the chemical abundances of Boötes I show discrepancies in elemental abundances between different authors,
and thus a consistent picture of its chemical enrichment history has not yet been established. In the present work, we independently
determine chemical compositions of six red giant stars in Boötes I, some of which overlap with those analyzed in the previous studies.
Based on the derived abundances, we re-examine trends and scatters in elemental abundances and make comparisons with MW field
halo stars and other dwarf spheroidal galaxies in the MW.
Methods. High-resolution spectra of a sample of stars were obtained with the High Dispersion Spectrograph mounted on the Subaru
Telescope. Abundances of 12 elements, including C, Na, α, Fe-peak, and neutron capture elements, were determined for the sample
stars. The abundance results were compared to those in field MW halo stars previously obtained using an abundance analysis techniquesimilar to the present study.
Results. We confirm the low metallicity of Boo-094 ([Fe/H]= −3.4). Except for this star, the abundance ratios ([X/Fe]) of elements
lighter than zinc are generally homogeneous with small scatter around the mean values in the metallicities spanned by the other fivestars (−2.7 <[Fe/H]< −1.8). Specifically, all of the sample stars with [Fe/H]> −2.7 show no significant enhancement of carbon. The[Mg/Fe] and [Ca/Fe] ratios are almost constant with a modest decreasing trend with increasing [Fe/H] and are slightly lower than thefield halo stars. The [Sr/Fe] and [Sr/Ba] ratios also tend to be lower in the Boötes I stars than in the halo stars.
Conclusions. Our results of small scatter in the [X/Fe] ratios for elements lighter than zinc suggest that these abundances were
homogeneous among the ejecta of prior generation(s) of stars in this galaxy. The lower mean [Mg/Fe] and [Ca/Fe] ratios relative
to the field halo stars and the similarity in these abundance ratios with some of the more luminous dwarf spheroidal galaxies atmetallicities [Fe/H]< −2 can be interpreted as star formation in Boötes I having lasted at least until Type Ia supernovae started tocontribute to the chemical enrichment in this galaxy.
Key words. nuclear reactions, nucleosynthesis, abundances – galaxies: abundances – galaxies: dwarf – galaxies: individual(BoötesI) – stars: abundances

1. Introduction
Ultra-faint dwarf spheroidal galaxies (UFDs), recently discovered in the Local Group, provide us with an excellent laboratory
for studying the formation of low-mass galaxies and production of heavy elements in their earliest stages in the Universe.
The faint Boötes I dwarf spheroidal galaxy (dSph) is one of themost interesting objects of its kind as a candidate for first galaxies and/or as building blocks of the Milky Way (MW). It was
first discovered with the help of photometric data obtained by
the Sloan Digital Sky Survey (SDSS; York et al. 2000) as oneof the faintest galaxies known to date (Belokurov et al. 2006).
Subsequent spectroscopy to measure line-of-sight velocities ofcandidate member stars suggests that Boötes I is a dark-matterdominated system with a total mass on the order of 107 M(Muñoz et al. 2006; Martin et al. 2007). Koposov et al. (2011)

has developed a reduction technique to derive more accurate radial velocities, and they suggest that the velocity dispersion ofthis system is represented better by two components with 2.4
km s−1 for a "cold" component and ∼ 9 km s−1 for a "hot" one,which could be interpreted as the result of velocity anisotropy inthis galaxy.
Detailed analyses of a color-magnitude diagram (CMD) for
Boötes I suggests that it consists of a very old stellar population with age > 10 Gyr, which is similar to the most metal-poor
globular clusters in the MW (Okamoto et al. 2012). Chemical
abundance analyses based on lower resolution spectroscopy report very low metallicities for Boötes I stars at [Fe/H]∼ −2.5 on
average with a large spread of ∆[Fe/H]& 1.0 dex (Norris et al.
2008; Lai et al. 2011).
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A&A proofs: manuscript no. boo16

Detailed chemical abundances for Boötes I allow us to further address whether possible progenitors of the metal-poorcomponent of the MW are similar to Boötes I and whether
this galaxy is indeed an example of the first galaxies formed in
the early Universe. However, these questions have not yet beenclearly answered because of the discrepancy in abundance results among previous studies. Chemical abundances based onhigh-resolution spectroscopy for Boötes I stars have previously
been obtained by Feltzing et al. (2009, hereafter F09), Norris
et al. (2010b), and Gilmore et al. (2013, hereafter G13). Feltzing
et al. (2009) analyzed seven stars with −3 .[Fe/H]. −2 basedon the high-resolution spectra using an LTE abundance analysis
code that takes the sphericity of stellar atmospheres into account.
They report the presence of a star with an anomalous [Mg/Ca]
ratio, Boo-127, and interpret it as a signature of nucleosynthesis
yields of individual supernove (SNe).
Gilmore et al. (2013) carried out the double-blind analysis by
two groups of the co-authors, Norris & Yong (NY) and Geisler& Monaco (GM), for six stars in Boötes I with −3.7 <[Fe/H]<−1.9. They find that the chemical abundance pattern for the
Boötes I stars are generally similar to those observed in field halostars with small scatter. However, in contrast to what is generally
known for the field halo stars with the similar metallicities, a signature of decreasing [α/Fe] ratios with increasing [Fe/H] is alsosuggested. They observed Boo-127 in common with F09 but did
not confirm the [Mg/Ca] anomaly of this object. Although thesedata provide important implications for the early chemical enrichment of this galaxy, they provide the contrasting pictures,
namely, the former study reports the presence of an anomalous
star, which may suggest that the chemical evolution in Boötes
I proceeded in an inhomogeneous manner, while the latter does
not support the presence of this inhomogeneity.
To obtain a consistent picture of the early chemical evolutionof Boötes I, we present an independent analysis of six Boötes I
stars with −3.2 <[Fe/H]< −2.0 based on high-resolution spectroscopic data obtained with the Subaru/High Dispersion Spectrograph (HDS). Five stars and three stars in our sample havebeen analyzed in common with F09 and G13, respectively.
This paper is organized as follows. In Section 2, we describe
our sample, the method of observation, and equivalent width(EW) measurements. Section 3 presents the method for the stellar parameter and abundance estimates for the sample stars. In
Section 4 we report the results of our abundance analysis andtheir comparison with previous results in the literature. Section
5 first presents the abundance comparison of the present Boötes I
sample with the field MW halo stars and with other dSphs. Then,possible implications of the chemical evolution of Boötes I are
discussed. Finally, Section 6 concludes the paper.

2. Observation and measurements
2.1. Sample and observation

The sample of six bright red giant branch stars in Boötes I was
selected from Norris et al. (2008) to cover a wide metallicity
range (−3 .[Fe/H]. −2). The high-resolution spectra of six
bright red giants were obtained with the Subaru Telescope High
Dispersion Spectrograph (HDS; Noguchi et al. 2002) on 16-17
May 2009 and 16 May 2010. The objects are listed in Table 1
with observation details. We applied CCD on-chip binning (2×2pixels), resulting in a resolving power of R = λ/δλ = 40, 000.
Our spectral ranges cover 4030–6800 Å. Spectra of two comparison stars in the Galactic halo (HD 216143 and HD 85773) were
Article number, page 2 of 15page.15

Table 2. The Mg linelists.

F09a
λ
(Å)
4571.104702.994730.035172.68
5183.60
5528.40
5711.09

χ
(eV)
0.00
4.35
4.35
2.71
2.72
4.35
4.35

...
−0.666...
−0.402−0.180−0.620...

G13(NY)a G13(GM)alog g f
(dex)
...
...
...
...
...
...
...
...
...
...
−0.340−0.620...
−1.833TWa−5.623−0.440−2.347...
...
−0.498...

Notes. (a) The log g f values for the Mg lines adopted in F09, G13, andthis work (TW). The two independent analyses carried out by G13 arereferred to as "NY" and "GM", respectively (see Section 1).

obtained with similar setups of the instrument in our previousrun (Aoki et al. 2009).
A standard data reduction was made with IRAF echelle package. Cosmic-ray hits were removed by the method described in
Aoki et al. (2005). The observations were carried out mostly
during dark nights, and the sky background was not significant.
One exposure of Boo–121 on 16 May 2010 was exceptionallyaffected by the Moon. We thus subtracted the sky backgroundestimated from the slit image around the object. The spectra obtained from individual exposures were combined to obtain final
spectra.
Radial velocities of the sample stars are measured from 2550 Fe I lines for each object. The results are given in Table 1(vhelio ). The measured values agree with those derived by G13
within 1.5 km s−1 for the three stars observed in common.
2.2. Equivalent width measurement

The EWs of absorption lines were measured by Gaussian fitting. The atomic line list used in Ishigaki et al. (2012, 2013) was
adopted. The errors in the EW measurements estimated from theformula presented in Cayrel (1988) range from ∼ 5 to 12 m Å,
depending on the signal-to-noise ratios (∼ 10 per pixel at ∼ 4500
Å and ∼ 30 at ∼ 6000 Å) of the continuum.
Although Mg is a key element in our work, the gf values
could be uncertain because of the lack of laboratory measurements for most of the Mg I lines (Aldenius et al. 2007). We inspected the gf values of Mg I lines adopted in previous studies
for Boötes I stars F09 and G13 and summarize them in Table 2(the values of F09 was provided by Feltzing, private communication). The last column gives the log g f values adopted in thiswork. The Mg I lines at 5172.68, 5183.60, and 5711.09 Å were
not used in this work because EWs of these lines for our sample
stars were too strong or too weak to be useful in the abundance
analysis. As can be seen from Table 2, the log g f values for the5528.40 Å line, which were common to all of these studies, aredifferent by up to ∼ 0.3 dex.
We confirmed that the Mg abundances derived from different lines using the log g f values adopted in this work (last column of Table 2) agree within ∼0.3 dex for the comparison stars,
HD 216143 and HD 85773, for each.
2.3. Equivalent width comparisons

Figure 1 presents the comparison between the EWs measured inthis work and those in G13 for the three stars analyzed in com-

M. N. Ishigaki et al.: Chemical compositions of six metal-poor stars in the ultra-faint dwarf spheroidal galaxy Boötes I
Table 1. Objects and observations

Name
Boo–094
Boo–121
Boo–911
Boo–127
Boo–009
Boo–117

RA(J2000)
14:00:31.50
14:00:36.52
14:00:01.07
14:00:14.56
13:59:48.81
14:00:10.49

DEC(J2000)
14:34:03.6
14:39:27.3
14:36:51.5
14:35:52.7
14:19:42.9
14:31:45.5

mon. For all of these stars, the EWs measured in this work aresystematically lower than the values obtained by G13. The meandifferences between the EW in this work and in the analysis of"NY" in G13 are 16-26 mÅ, while those between in this work
and in the other independent analysis, "GM", in G13 are 10-18
mÅ, which are modestly larger than the errors in the EWs measurement (∼ 12 mÅ).
To test whether the uncertainty in the continuum placement
could be the reason for the discrepancy, we examined the change
in the measured EWs by adopting different continuum levels forone of the Fe I line as shown in Figure 2. For this line, the discrepancy in the measured EWs among the three analyses is particularly large (72.0 mÅ, 96.2 mÅ and 102.5 mÅ in this work,
GM, and NY, respectively). If the continuum level of the normalized spectrum is changed by ±0.05, which is slightly morethan the scatter in the continuum around this line (widths in theshaded regions), the measured EW varies in the range 52.7−89.2
mÅ. This result illustrates that the variation in the continuum
level partly explains the EW difference among the three analyses. However, as can be seen in Figure 2, the error in the continuum level that is as large as ±0.05 is unlikely for this line, so it
is difficult to completely illuminate the discrepancy. In this paper, we use the measured EWs without any correction. Instead,
we examine the effects of this offset in derived abundances in
Section 3.3.

3. Abundance analysis
3.1. Stellar atmospheric parameters

Effective temperatures (T eff ) were estimated based on griz color
indices using the temperature scale of Castelli1 . The griz photometry was adopted from the SDSS DR9 data set. The Galactic extinction in the griz magnitudes toward individual stars
was corrected based on the revised extinction measurement by
Schlafly & Finkbeiner (2011) via the NASA/IPAC Extragalactic Database2 . The extinction in the g-band varies from 0.052 to
0.060 among our six sample stars. These values are translated
into E(B − V) in the Schlegel et al. (1998) from 0.016 to 0.018(Schlafly & Finkbeiner 2011), which are consistent with the previously adopted value of E(B − V)= 0.02 (Norris et al. 2008).
The surface gravity (log g) for the sample stars are then estimated by comparing the adopted temperature with a 12 GyrYonsei-Yale (YY) isochrone (Demarque et al. 2004) by assuming that the stars lie on the red giant branch of the system. In
the above procedure for obtaining T eff and log g, the [Fe/H] values are needed as input. We adopt the initial guess of the [Fe/H]
values from Norris et al. (2010a).
The microturbulent velocity (ξ) is estimated by requiring that
the Fe abundances from individual Fe I lines do not depend on12

http://wwwuser.oat.ts.astro.it/castelli/colors/sloan.html
http://ned.ipac.caltech.edu/forms/calculator.html

g
17.50
17.92
17.97
18.16
17.93
18.20

(g − r)0
0.87
0.82
0.80
0.77
0.81
0.74

vhelio (km s−1 )95.0±0.7
105.4±0.6102.2±0.6
98.5±0.597.1±0.4
98.0±0.6

HJD
2454968245496824549692455333
2455333
2455333

Fig. 2. EW measurements for one of the Fe I lines by adopting the different continuum levels. The dotted line shows the best-fit Gaussian,where continuum level is estimated from the shaded wavelengths region. The dashed lines show the Gaussian fit by fixing the continuum
levels to be ±0.05 of the best-fit value.

their measured EWs. In this step, we have used the Fe I lines
having a reduced EW log(EW/λ) < −4.7.
After the standard 1D-LTE abundance analysis is performed,
the procedures are iterated until a consistent set of parameters(T eff , log g, ξ, and [Fe/H]) is obtained. The resulting atmosphericparameters are summarized in Table 3.
The adopted T eff and log g in this work and in the two independent analyses (“NY” and "GM") presented by G13 agree
within 100 K and 0.2 dex, respectively, for the three stars analyzed in common. The ξ values are different between the twoanalyses in G13 by 0.5-1.3 km s−1 . The relatively high ξ value
for Boo–094 reported by the "NY" analysis is in excellent agreement with ours, while the ξ values for Boo–117 and Boo–127
are more consistent with those obtained by the "GM" analysis.
3.2. Abundances

Abundances of individual elements are determined by standard
analyses using the measured EWs. We restrict the analysis tothe absorption lines with log(EW/λ) < −4.7 to minimize theerrors in derived abundances. For Na I lines, we relax this restriction to log(EW/λ) ≤ −4.5, since only strong resonance lines
are detected for this species. Similarly, only strong resonance Sr
II lines are available for the abundance estimate of Sr. A specArticle number, page 3 of 15page.15

A&A proofs: manuscript no. boo16

Fig. 1. Comparison of the measured equivalent widths (EWs) in this work (TW) and in the two analyses, "NY" (closed circles) and "GM" (opencircles), in Gilmore et al. (2013).
Table 3. Atmospheric parameters

ObjectBoo–094
Boo–121
Boo–911
Boo–127
Boo–009
Boo–117
HD216143
HD85773

T eff
(K)
4500
4500
4500
4750
4750
4750
4529
4366

log g
(dex)
0.8
0.8
0.9
1.6
1.4
1.5
1.3
1.0

ξ
(km s−1 )
3.4
2.1
1.7
1.6
2.5
1.9
1.9
2.1

[Fe/H]
(dex)
−3.2
−2.5
−2.2
−1.9
−2.7
−2.2
−2.1
−2.4

tral synthesis is applied to estimate Sr abundances as describedin detail in Section 4.7. The hyperfine splitting effect is includedin the analysis of Ba lines (McWilliam 1998), assuming the rprocess isotope ratios. The EWs and abundances estimated fromindividual lines are summarized in Table 4.
Carbon abundances are determined from the comparisons ofsynthetic spectra to the observed spectrum of the CH molecularband at around 4323 Å. For Boo–094, Boo–121, and Boo–009,
only an upper limit of carbon abundance is estimated. Because
the signal-to-noise ratios are low for the wavelength range of
Article number, page 4 of 15page.15

the G-band, the continuum level is not determined well, which
produces the abundance uncertainty of 0.1 − 0.3 dex.
3.3. Abundance errors

The errors due to the line-to-line scatter (σline ), which are calculated as the standard deviation of the abundances from individual
lines divided by the square root of the number of lines, are summarized in Table 5. If only one line is available for the abundance
estimate, σline is assumed to be equal to the standard deviation
for the Fe I lines, which are typically ∼ 0.2 dex in the present
work.
The abundance errors due to uncertainties in T eff , log g, andξ are examined by calculating the abundance differences whenchanging these parameters by ±100K, ±0.3 dex, and ±0.3 km
s−1 , respectively. Results of this exercise are summarized in Table 5. The error bars in Figures 5 and 6, and σtot in Table 5 correspond to a quadratic sum of the contribution from the uncertainty
in the three atmospheric parameters and line-to-line scatter in theabundances.
The systematic offset in the EWs from those in Gilmore
et al. (2013) mentioned in Section 2.3 are an additional sourceof uncertainty in measuring absolute abundances in this work.

M. N. Ishigaki et al.: Chemical compositions of six metal-poor stars in the ultra-faint dwarf spheroidal galaxy Boötes I

Fig. 3. Changes in abundances when all of the EWs in this work areincreased by 15mÅ to partly correct for the EWs offset from Gilmore
et al. (2013).

F09. This difference can be partly attributed to the higher log g
and larger ξ adopted in our analysis.
The abundances of Mg and Ca also marginally agree withthose estimated by F09 within 0.30 dex, except for Boo-127.
For this star, F09 report an exceptionally high [Mg/Ca] ratio([Mg/Ca]> 0.6), while this work has obtained [Mg/Ca]=0.08.
The estimated log (Mg) in this work is ∼ 0.5 dex lower thanthat reported by F09. Several factors seem to be responsible forthe abundance difference, including (1) a higher log g f value for
Mg I 5528 Å (see Table 2), (2) higher log g, and (3) higher ξ inthis work than those adopted in F09. While the error in the EWs
by up to ∼ 10 − 15 mÅ may also contribute to the discrepancy, it
does not totally explain the large difference.
To illuminate the uncertainty in the log g f values, we also
performed a simple differential analysis to obtain [Mg/Ca] for
Boo–127 with respect to one of the comparison stars, HD
217143. Since the T eff , log g, and [Fe/H] values are similar between the two stars, a direct comparison of EWs for the Mg I
and Ca I lines provides the [Mg/Ca] estimate. Figure 4 shows
the wavelength regions that contain the Mg I and Ca I lines for
Boo-127 and HD 216143. While both stars show Ca I line withsimilar strengths, the Boo-127 shows slightly weaker Mg I line
than HD 216143, which does not support the anomalously high
[Mg/Ca] in Boo-127 compared to the reference star. More quantitatively, The relative abundance log(A/Aref ) is expressed as
!
!
!
A
EW
κν
= log
− log
− ∆θex χAref
EWrefκν,refFigure 3 shows changes in log (X) abundances when the EWs
are increased by 15mÅ to reduce the systematic offset from theGilmore et al. (2013) analyses for the most metal-poor (Boo–094) and metal-rich (Boo–127) stars in our sample. As can beseen, after the EW offset is applied, the abundances are systematically increased by ∼ 0.15 − 0.30 dex. On the other hand, theamount of correction is mostly similar for all elements, and thus
difference in the abundance ratios relative to iron ([X/Fe]) partly
cancel out. Although the absolute abundances could be affectedby the systematic errors in EW measurements, the abundance
ratios ([X/Fe]) are consistent with those in Gilmore et al. (2013).
For the Sr abundances estimated by the spectral syntheses,
uncertainty in the continuum level (1σ ∼ 0.1) causes an error aslarge as ±1.0 dex. as detailed in Section 4.7.

log

4. Abundance results

4.2. Comparison with Gilmore et al. (2013)

Our abundance results are summarized in Table 5. Relative abundances ([X/Fe]) were determined adopting the solar abundancesfrom Asplund et al. (2009). Figures 5 and 6 plot the relative
abundances of C, Na, Mg, Ca, Sc, Ti I, Ti II, Cr, Ni, Zn, Sr, andBa as a function of [Fe/H] for the sample of six Boötes I stars(filled circles) and the two comparison field halo stars (open circles). Triangles represent the abundance results from the "GM"analysis in G13, whose systematic offset in the EWs from thosein this work is smaller than that in the “NY" analysis.

Comparisons of the relative abundances in this work and in the"GM" analysis of G13 are presented in Figure 5 (triangles thatare connected with the data points in this work) for the three stars
analyzed in common. Within the quoted error bars, our abundance results generally agree with those obtained by G13 witha few exceptions. One relatively large discrepancy is found in
the [Fe/H] value for Boo-094, the most metal-poor star in our
sample, which is ∼ 0.2 dex lower in this work than obtained by
G13. The smaller EWs (see Section 2.3) by ∼ 10mÅ and larger
ξ adopted in this work than in G13 are probably responsible forthe discrepancy.

4.1. Comparison with Feltzing et al. (2009)

Five objects in the present sample have also been studied by F09,
who determined Fe, Mg, Ca and Ba abundances for seven objects
in Boötes I. The [Fe/H] of the sample stars determined by thepresent work agree well with the results of F09 within ∼ 0.1 dex,
except for the lowest metallicity star Boo-094 in our sample, forwhich the estimated [Fe/H] is −3.4 dex in our study but −2.9 in(1)

where κν and κν,ref are continuum opacity, ∆θex (θ = 5040/T )is the difference in an excitation temperature, and χ is the excitation potential of the line (Gray 2005). If we assume that theκν term is negligible and ∆θex is constant and equal to the one
computed from the effective temperature (∆θex = ∆θex,eff ) for all
relevant optical depths, the relative abundance can be obtained
from the ratios of the EWs (Smith & Harmer 1982). This resultsin the [Mg/Ca] ratio for Boo–127 being ∼ 0.04 dex lower thanthat of HD 216143, which is, again, inconsistent with the anomalously high [Mg/Ca] ratio. The low, but not extremely low, Ba
abundances ([Ba/Fe]∼ −0.6) of Boo-121 and Boo-911 found by
F09 are confirmed by our measurement. The log (Ba) in the twostudies agree within 0.3 dex.

4.3. Carbon

The carbon abundance is estimated for three of the six sample
stars, while only an upper limit is obtained for the other threeobjects. For the six sample stars, extremely carbon-rich objects
are not found. These sample stars all show [C/H]< −2.7 andArticle number, page 5 of 15page.15

A&A proofs: manuscript no. boo16

4.6. Iron-peak elements

Fig. 4. Spectra of Boo–127(red) and the comparison star, HD 216143(black), for the wavelength regions that contain Mg I and Ca I lines.

[C/Fe]< 0.2. For the three objects with measured carbon abundances, the average [C/Fe] is −0.7 dex with small scatter. These
values are close to the lowest bounds of the [C/Fe] distributions
obtained by Lai et al. (2011) and Norris et al. (2010a) for thisgalaxy. This [C/Fe] value is comparable to what is seen in field
halo stars having similar luminosity (Spite et al. 2005; Norris
et al. 2010a).
4.4. Sodium

For the present Boötes I sample stars, sodium abundances
are only estimated using the strong Na I resonance lines at
5889/5895 Å. These lines are known to suffer from strong
non-LTE (NLTE) effects such that LTE abundances tend to belarger than corresponding NLTE values (Takeda et al. 2003). The
NLTE correction to Na abundances ranges from −0.1 dex to −0.5
dex, depending on the atmospheric parameters of stars (Takeda
et al. 2003). Without any NLTE correction, the average [Na/Fe]
value for the six Boötes I stars is ∼ −0.1 dex with small scatter.
This value is similar to the uncorrected [Na/Fe] values obtainedfor field halo stars with comparable metallicities.
4.5. α/Fe ratios

The [Mg/Fe] ratios in our sample stars are similar to or slightly
lower than those obtained for the field halo stars. A hint of amodest decreasing trend of the [Mg/Fe] with increasing [Fe/H]
is seen, as was also reported in G13. The [Ca/Fe] ratios are alsosimilar to or slightly lower than the value for the field halo stars.
Excluding the lowest metallicity star, a scatter in the [Ca/Fe] ratios is remarkably small (∼ 0.1 dex). The [Ti I/Fe I] and [Ti II/Fe
II] ratios differ by up to ∼ 0.3 dex in our Boötes I sample stars,
as reported for the field halo stars (e.g., Lai et al. 2008). The
Boötes I stars tend to show slightly lower [Ti I/Fe small I] abundance ratios than the halo stars, while the [Ti II/Fe II] ratios forthese samples are similar.
Article number, page 6 of 15page.15

Scandium abundances are estimated from a few to several Sc II
lines in our sample stars. The mean abundances are similar toor slightly lower than those reported for the field halo stars in[Fe/H]> −2.5.
Chromium abundances are obtained from a few Cr I lines for
the Boötes I stars. The [Cr/Fe] ratios are similar to those reported
for the field halo stars for two of the most metal-rich stars, whileother stars show lower [Cr/Fe] ratios. However, since only one
or two lines with low-signal-to-noise ratio are used in the abundance estimate, the significantly lower [Cr/Fe] ratios for thesestars require further confirmation.
In the four higher metallicity stars in our sample, the [Ni/Fe]
ratios are similar to those obtained for the field halo stars, except for one object (Boo–121). Two Ni I lines used to obtain Ni
abundance in this stars both suggest a subsolar [Ni/Fe] ratio.
Since EWs of Zn I lines are not reliably measured for oursample stars, we used spectral synthesis to estimate upper limitsof the Zn abundances. The results are shown by bars with downward arrows in the bottom right hand panel of Figure 5. The
upper limits of the [Zn/Fe] ratios are similar to or lower than thefield halo samples in the metallicity range [Fe/H]> −2.5.
4.7. Neutron-capture elements

The strontium abundances are obtained by fitting synthetic spectra to the observed ones for the Sr II resonance line at 4215.5
Å. The log g f value for this line is taken from NIST database.
Figure 7 shows the resulting synthetic and the observed spectra
around the Sr II line for the Boötes I stars for which the Sr abundances are reasonably obtained. The other three sample stars are
excluded in the analysis since the spectral region is too noisy toestimate Sr abundances.
In the fitting procedure, we redefine the continuum level using the wavelength regions 4214.10–4215.25 Å and 4216.80–
4217.10 Å. To estimate abundance errors due to the uncertainty
in the continuum level, we performed the fitting changing thecontinuum level of the synthetic spectrum by ±1σSTDV , where
σSTDV is a standard deviation in the observed continuum level.
The error estimated in this procedure is ∼ 0.7 − 1.5 dex, as summarized in Table 5.
For comparison, the bottom panel of Figure 7 shows the observed spectrum of a field halo giant, HD 107752, which has
atmospheric parameters similar to those of the present sample(T eff = 4820 K, log g= 1.6, ξ = 1.9 km s−1 , and [Fe/H]=−2.8)(Ishigaki et al. 2013). While the Sr II line of the HD 107752
appears to be stronger than a nearby Fe I line, those in the three
Boötes I stars seem to have comparable strength as the Fe I lines.
Andrievsky et al. (2011) studied the NLTE effects on Srabundance determination for extremely metal-poor stars and report that the NLTE correction for the Sr II 4215.5 Å line is at
most ∼0.2 dex for stars having atmospheric parameters similar
to the present Boötes I sample. The correction is smaller than theother sources of uncertainty in the present analysis, so that is notincluded in our abundance results.
The top panel of Figure 6 shows resulting [Sr/Fe] values plotted against [Fe/H] with the error bars. The crosses and small
filled circles represent the abundances from literature extracted
from SAGA database and those from Ishigaki et al. (2013), respectively. As can be seen, even though the errors are very large,
the [Sr/Fe] ratios of the three Boötes I stars are significantly
lower than the solar value, which is not typical for the field halostars at similar metallicities. Most of the field halo sample inM. N. Ishigaki et al.: Chemical compositions of six metal-poor stars in the ultra-faint dwarf spheroidal galaxy Boötes I

Fig. 5. [X/Fe] ratios plotted against [Fe/H] for the sample of six Boötes I stars (red circles) and the two comparison field halo stars (blue circles).
Black triangles represent the abundance results from the "GM" analysis in G13. Small gray circles present the abundances of field halo stars from
Ishigaki et al. (2012, 2013). In the plot for [C/Fe]-[Fe/H], results from Norris et al. (2010a) for Boötes I stars are shown in crosses.

literature show the solar value in [Fe/H]< −2.5 and increasing dispersion toward lower metallicity with [Sr/Fe]∼ −2.0 to∼ 0.5. On the other hand, the Boötes I stars have [Sr/Fe]∼ −1 in[Fe/H]> −2.5, and the most metal-poor star in our sample showeven lower [Sr/Fe] of ∼ −2.
The bottom panel of Figure 6 plots the [Ba/Fe] ratios against[Fe/H]. The [Ba/Fe] of Boötes I stars are subsolar, generally confirming the result of G13. The [Ba/Fe] ratios for these stars tend
to distribute at the lower bound of the [Ba/Fe] distribution for
the field halo stars.
One interesting feature is that Boo–127, which is the mostmetal-rich star in our sample, shows lower [Ba/Fe] than the field
halo stars at similar metallicities ([Fe/H]∼ −2). As shown in Fig
4, the Ba II line is significantly weaker in Boo–127 than seen in
the comparison star with similar [Fe/H].

5. Discussion
5.1. Comparison with the field MW halo

If the MW halo was assembled from smaller satellite galaxies
similar to the surviving dSphs, as suggested by the standard cosmological model, we naively expect that chemical abundancesin the MW field halo stars are similar to these dSphs. In thissection, we address the question of whether the observed abundances in Boötes I stars are similar to those obtained for the field
halo stars.
Figure 8 compares mean abundances of the Boötes I stars andfield halo stars in the metallicity range of −3.0 <[Fe/H]< −2.0.
To minimize systematic errors in the abundance comparison,eight halo giants (T eff < 5000K and log g< 3.0, including oneof the comparison stars HD 85773) from Ishigaki et al. (2012,2013) are only considered for the field halo sample. In theirwork, the adopted 1D-LTE abundance analysis technique andthe line list are the same as in this work. Error bars in Figure
8 represent standard deviations around the mean values.
In comparison to the field halo stars, the Boötes I stars show
slightly lower [Mg/Fe] and [Ca/Fe] ratios in the metallicity rangeof −3.0 <[Fe/H]< −2.0. Kolmogorov-Smirnov (KS) tests suggest that the probability at which [X/Fe] ratios of the halo stars
and the Boötes I stars can be drawn from the same distributionis less than 5% for Mg and Ca, which marginally supports theabundance difference. On the other hand, the [Na/Fe] ratios for
the Boötes I stars are indistinguishable from those of the halostars.
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A&A proofs: manuscript no. boo16

Fig. 6. Same as Figure 5 but for Sr and Ba. Abundance data extracted
from the SAGA database (Suda et al. 2008) (RGB, R > 30000 withno flags of "C-rich" nor “CEMP") are additionally plotted with graycrosses to supplement the MW halo data.

As mentioned in the previous sections, the Boötes I and the
MW halo stars show a hint of an abundance difference in Sr. Figure 9 shows log (Sr) plotted against log (Ba) for the Boötes I
stars and MW halo giants. The data for the MW halo giants wereextracted from the SAGA database (Suda et al. 2008) with criteria, −3.5 <[Fe/H]< −2.0 and a spectral resolution R > 30000
with no flags of “C-rich" or “CEMP”. As can be seen, most of
the halo stars show Sr/Ba> 1, while the Boötes I stars show
Sr/Ba. 1. Aoki et al. (2005) found that some of the halo stars
with [Fe/H]< −3.0 show Sr/Ba> 1, while the excess of Sr is notseen for stars with higher Ba abundances. They interpret the result as evidence that the two astrophysical sites are responsible
for producing these neutron-capture elements, one of which produces both Sr and Ba and the other produces more Sr than Ba.
Since the signature of enhanced Sr relative to Ba is not evident in
our Boötes I sample, the latter astrophysical site may be lacking
or less frequent in this galaxy. We discuss the origin of Sr andBa in Bootes I in Section 5.3.5.
To summarize, the present sample of Boötes I stars showa hint of difference in abundances of some of the α-elements
and neutron-capture elements from the bulk of MW halo stars.
This result suggests that the chemical abundances of Boötes I
stars alone are not enough to fully explain the observed chemical abundances for field halo stars.
Corlies et al. (2013) have investigated conditions that yielda difference in chemical abundances between surviving satellite
galaxies and progenitors of the MW halo using an N-body simulation. They suggest that the progenitors of the stellar halo usually experienced more cross pollution by neighboring subhalos,
while the surviving satellites tend to have evolved chemicallyin isolation. Whether the degree of the cross pollution explains
the observed abundance difference may have interesting implications for the hierarchical formation of the MW halo.
Since the present result is based on a small sample, a larger
sample is clearly desirable to extract a conclusion for the range
in elemental abundances spanned by the whole stellar population in Boötes I. Moreover, abundances of field halo stars areoften determined only for relatively nearby objects (within a several kpc), while the estimated distance to Boötes I is 64±3 kpc(Okamoto et al. 2012). A systematic abundance survey for moredistant halo stars would be needed to examine possible connection between dSphs similar to Boötes I and the MW field halo.
5.2. Comparison with other dSphs

Fig. 7. Observed (black dots) and the best-fit synthetic (solid line) spectra of the three Boötes I stars around the Sr II 4215.5 Å line. The obtained [Fe/H] and [Sr/H] values are indicated in the bottom-left cornerof the each panel ("Name/[Fe/H]/[Sr/H]"). The dotted lines show thesynthetic spectra fitted after changing the continuum levels by ±1σSTDV(in each panel, the continuum levels have been shifted to match the bestfit value). For comparison, the spectrum of one halo giant, HD 107752,
is shown in the bottom panel.
Article number, page 8 of 15page.15

Figure 10 compares the [Mg/Fe] ratios of the Boötes I stars inthis work and those of other MW dSphs that include very metalpoor ([Fe/H]< −2) stars from the literature (Aoki et al. 2009; Cohen & Huang 2009, 2010; Frebel et al. 2010; Koch et al. 2008).
Relatively luminous (LV, & 105 L ), so-called classical dSphs(Draco, Sextans, Ursa Minor) are shown with green symbols,
while less luminous dSphs (UFDs; Boötes I, Coma Berenices,
Hercules, Ursa Major II) are shown in red symbols. According
to these available data, the more luminous dSphs appear to show
relatively low [Mg/Fe] ratios than field halo stars at [Fe/H]∼ −2
and/or a decreasing [Mg/Fe] trend with increasing [Fe/H]. On
the other hand, the fainter dSphs such as Coma Berenices or Ursa
Major II appear to show comparable or higher abundance ratioswith large scatter. The Boötes I stars generally show the [Mg/Fe]
ratios more similar to some of the classical dSphs, such as Ursa
Minor rather than the fainter UFDs, such as Coma Berenices or
Ursa Major II, although there seems to be large variation in theabundance ratios among the classical dSphs.
To examine whether the abundance ratios can be related to
the global properties of dSphs (e.g., dynamical masses or luminosity), the lower panels in Figure 10 plot average [Mg/Fe] ratiosof these dSphs against their central line-of-sight velocity dispersions (σlos ; left) and logarithms of total luminosity (L; right).

M. N. Ishigaki et al.: Chemical compositions of six metal-poor stars in the ultra-faint dwarf spheroidal galaxy Boötes I

Fig. 8. The abundance pattern of the Boötes I stars (filled circles) andthe MW halo stars (open circles) with −3.0 <[Fe/H]< −2.0.

The exceptionally high [Mg/Fe] ratios are reported for the two
Hercules stars, which may suggest an anomalous star formation
history for this galaxy (see also François et al. 2012). According
to the available data, differences in other elemental abundances
lighter than zinc, such as [Na/Fe] ratios, are not very clear among
all dSphs including Boötes I for [Fe/H]< −2.
The MW dSphs have been reported to follow a correlation
between mean [Fe/H] and L (Kirby et al. 2011b) extending fora wide range of luminosities (3.5 < log(L/L ) < 7.5). Although
precision and sample size may not be sufficient to examine the[Mg,Ca/Fe]-σ, -L relation, if any, from currently available data,
the locations of Boötes I in these diagrams may have implications for the chemical evolution of this galaxy, which we discussin Section 5.3
Figure 12 shows [Sr/H] and [Ba/H] ratios plotted against[Fe/H] for Boötes I and other dSphs. Classical dSphs generally show increasing [Sr/H] and [Ba/H] ratios with increasing [Fe/H], similar to the trend seen for field halo stars. On
the other hand, the abundances for fainter dSphs tend to distribute at lower bounds than for the field halo stars. Lower abundances of neutron-capture elements for very-metal-poor stars in
dSphs have already been noticed in previous studies (e.g., Fulbright et al. 2004; Koch et al. 2008; Frebel et al. 2010; Honda
et al. 2011; Koch et al. 2013). Our sample of Boötes I stars
show [Ba/H] ratios that are more similar to the classical dSphs,
while [Sr/H] ratios are lower, similar to those observed in fainterdSphs.
Whether the neutron-capture elements are more deficientin dSphs than in the MW field halo stars needs to be investigated with a larger sample of high-quality spectra for distantobjects, which may be challenging with the current instrumentation. Nevertheless, the detection of Sr and Ba in Boötes I suggests that at least one event that produces neutron-capture elements must have occurred at the earliest epoch of star formationin this galaxy (cf., Roederer 2013).
5.3. Chemical evolution of Boötes I

Fig. 9. log (Ba) versus log (Sr) for the three Boötes I stars. For comparison, the abundance data for MW halo giants are extracted fromSAGA database (Suda et al. 2008) with the criteria −3.5 <[Fe/H]< −2.0
and a spectral resolution R > 30000 with no flags of “C-rich" or“CEMP”.

The data for σ and L are taken from Walker et al. (2009) except
for the σ of Boötes I, for which an updated value from Koposov+0.8
et al. (2011) (4.6−0.6
km s−1 , assuming a single component velocity distribution) is available. These figures suggest that thedSphs with larger velocity dispersions or higher luminosity tend
to show lower [Mg/Fe] ratios. Again, Boötes I, which has σ as
low as that of Coma Berenices and L midway between these twoclasses of dSphs, show the mean [Mg/Fe] ratio to be more similar to the classical dSphs, such as Ursa Minor or Draco. A similar trend can be seen for [Ca/Fe] ratios as shown in Figure 11.

In Section 5.1, we reported a possible offset in the abundancesof α and neutron capture elements between the Boötes I stars
and the field halo stars in the metallicity range −3 <[Fe/H]< −2.
These differences, if confirmed with larger samples, may suggest the differences in chemical evolution history between the
Boötes I and dominant progenitors of the MW stellar halo. As
discussed in Section 5.2, the [Mg/Fe] and [Ca/Fe] ratios in theobserved Boötes I stars are generally similar to the values reported for brighter (classical) dSphs in the metallicity range−3 <[Fe/H]< −2, while their velocity dispersion and luminosities are more similar to those of UFDs (Figures 10 and 11).
In this section, we discuss possible implications of the resultsof these comparisons on the chemical evolution of Boötes I in
terms of 1) contribution from Type Ia SNe, 2) galactic winds,
3) tidal interaction with the MW, 4) stochastic effects, and 5)
individual supernova yields.
5.3.1. Contribution from Type Ia SNe

In chemical evolution models, trends in abundances of αelements with [Fe/H] are diagnostics of the Type Ia SNe contribution relative to Type II SNe to chemical enrichments in galaxies (e.g., Matteucci & Greggio 1986). For many of known dSphs,
decreasing trends of [α/Fe] with increasing [Fe/H] are reported(Tolstoy et al. 2009; Kirby et al. 2011a; Vargas et al. 2013; Kirby
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A&A proofs: manuscript no. boo16

Fig. 10. Top panel: Abundance comparisons with classical (Draco: triangles, Sextans: three pointed asterisk, Ursa Minor: square) and ultrafaint (Coma Berenices: pentagon, Hercules: star, Ursa Major II: asterisk) dwarf galaxies from the literature (Aoki et al. 2009; Cohen &Huang 2009, 2010; Frebel et al. 2010; Koch et al. 2008). For comparison, the abundance data for MW halo giants extracted from SAGA
database (Suda et al. 2008) (R > 30000 with no flags of “C-rich" or“CEMP”) and from Ishigaki et al. (2012) are plotted with crosses anddots, respectively. Bottom panels: Mean [Mg/Fe] ratios for each galaxyplotted against velocity dispersion (σ, left) and luminosity (log L(L ),right) for that galaxy. The σ and L data are taken from Walker et al.
(2009).

et al. 2013), which could be indicative of chemical enrichment
proceeded partly through Type Ia SNe.
Whether Boötes I has experienced a significant contribution from Type Ia SNe to its chemical enrichments is an important question for constraining the star formation history inthis galaxy. Photometric observations suggest that Boötes I isdominated by an old stellar population with little age spread(Okamoto et al. 2012), which suggests that the duration of starformation was shorter than the uncertainty of the photometricage estimates (a few Gyr). Gilmore et al. (2013) suggest a contribution of Type Ia SNe in order to interpret the apparent decreasing [α/Fe] trend with [Fe/H]. On the other hand, the lowresolution spectroscopic study of Lai et al. (2011) does not finda signature of decreasing [α/Fe] ratios with increasing [Fe/H],
which does not support the Type Ia contribution.
The observed [Mg/Fe] and [Ca/Fe] ratios for the Boötes I
stars in this work appear to show a modest decreasing trend with
[Fe/H]. The Pearson test for linear correlation indicates that theprobability for a null correlation is 30 % for [Mg/Fe]-[Fe/H] and7 % for [Ca/Fe]-[Fe/H], thus only marginally supporting the correlation. Although a firm conclusion about the trend could not be
obtained for the present sample, the hint of the decreasing trend
suggests that the star formation in Boötes I lasted at least until
Type Ia SNe began to contribute to the chemical enrichment.
The decreasing [Ca/Fe] trend with [Fe/H] has also been reported for a sample of red giant stars in Hercules dSph, whichspans the [Fe/H] range similar to the present Boötes I sample.
This may imply that the contribution of Type Ia SNe to the chemical enrichment is common among galaxies as faint as Hercules
Article number, page 10 of 15page.15

Fig. 11. Same as Figure 10 but for [Ca/Fe].

Fig. 12. Comparison of the [Sr/H] and [Ba/H] ratios in Boötes I stars
obtained in this work and those in very metal-poor stars in other dSphs.
Symbols are the same as in Figure 10.

or Boötes I, if the Type Ia SNe is indeed responsible for lowering [Ca/Fe]. Analyses of the delay-time distribution of Type Ia
SNe suggest that a certain fraction of Type Ia SNe occur within< 1 Gyr of star formation (Maoz et al. 2010), which may support
their contribution to the chemical enrichments in Boötes I.
An alternative interpretation of the lower [Mg, Ca/Fe] ratios in Boötes I is a different initial mass function (IMF) inthis galaxy in comparison to the dominant progenitors of the
MW halo. Since [Mg/Fe] yields in Type II SNe could dependon masses of the progenitor stars (Kobayashi et al. 2006), a difference in IMF may lead to systematic difference in these abundance ratios.

M. N. Ishigaki et al.: Chemical compositions of six metal-poor stars in the ultra-faint dwarf spheroidal galaxy Boötes I

5.3.2. Galactic wind

Galactic winds driven by SNe are thought to significantly affect subsequent chemical evolutions in dSphs (Arimoto & Yoshii
1987). Chemical evolution models of Lanfranchi & Matteucci(2004) particularly suggest that the strength and time of occurance of the galactic winds play an important role in determining
the [α/Fe] ratios in the dSphs.
The winds are expected to occur when part of the energy injected by SNe heat up interstellar gas so that thermal energy ofgas exceed the binding energy of the galaxy (Arimoto & Yoshii
1987). In this scenario, the strength of the wind may dependon both the initial gas content and gravitational potential of thegalaxy.
As can be seen in the lower left-hand panels of Figures 10
and 11, the dSphs with smaller velocity dispersion, and thus
lower dynamical masses, such as Coma Berenices or Ursa Major II, show higher [Mg/Fe] and [Ca/Fe] ratios than the dSphs
with larger velocity dispersions. The likely correlation between
the velocity dispersion and the abundance ratios suggests thatthe chemical abundance depends on the gravitational potentialof galaxies, if the velocity dispersion traces the dynamical massof these systems.
Using the measured velocity dispersion profile, Walker et al.
(2009) estimate the dynamical mass of Boötes I within the halflight radius of this galaxy (rhalf = 242 ± 21 pc) to be ∼ 106 M .
If we assume that the mass has not changed significantly over
the age of this galaxy and if duration of star formation is indeed
longer than the timescale of Type Ia SNe, as suggested in the previous section, then the above scenario implies that a dark matter
halo with mass as low as that of Boötes I could host enoughinitial gas to retain a certain duration of star formation, or thatgalactic winds were not strong enough to truncate all subsequent
star formation in such a halo. Whether this scenario could bepossible depends on the possible tidal effect on the initial gas
and dark matter content of this galaxy as discussed below.
5.3.3. Tidal Interaction with the MW

We discussed in Section 5.2 a possible connection of chemical
abundance ratios with global properties, such as σlos and L. The
correlation is not currently clear, and Boötes I in particular appears to deviate from a possible relation. Thus, we may ask whatmakes Boötes I smaller σlos and lower L, while experiencing
chemical evolution that is more like the brighter classical dSphs.
One possible mechanism that may affect σlos or L is tidal interaction with the MW potential. Tidal interaction of Boötes Iwith the MW may have stripped the original content of baryons
and dark matter in this galaxy. N-body simulations for dynamical evolution of Boötes I suggest that this galaxy has lost asignificant amount of luminous and dark matter contents by thetidal forces as it orbits around the MW potential (Fellhauer et al.
2008). It is possible that Boötes I was initially as luminous as theclassical dSphs, while it has lost a significant amount of materials.
Since the full orbital motion of Boötes I is not known, it
is difficult to estimate how strong the tidal effects were in thisgalaxy. The tidal radius of this galaxy at its present location (62kpc from the Galactic center) is estimated to be 1.2 kpc (Fellhauer et al. 2008), which is larger than the luminous component
in this galaxy. Therefore, it is unlikely that the tidal effect is currently significant in stripping material from this galaxy. On theother hand, Collins et al. (2013) report that the estimated maximum circular velocity (Vmax ) for Boötes I is relatively low com-pared to other local group dSphs with wide ranges of luminosities, which could be explained by tidal stripping of dark matter. Therefore, one explanation of the observed similarity may
be that Boötes I once experienced similar chemical evolution
to, say, Ursa Minor, then within a few Gyr, it followed an orbitthat brought it closer to the Galactic center, experiencing strong
tidal force that were able to strip luminous and dark matter contents in this galaxy. All subsequent star formation stopped, which
results in the current low luminosity and very low metallicity([Fe/H]< −1.5) in this galaxy (Norris et al. 2010a). The information about the orbital motion of Boötes I would be necessaryfor examining such a scenario to interpret the observed chemical
abundances in this galaxy.
5.3.4. Stochastic effect

In very low-metallicity systems like Boötes I, it is possible that
only a small number of SNe of a prior generation of stars havecontributed to the chemical enrichment. For example, Frebel &Bromm (2012) suggest that first galaxies that formed exclusively
from ejecta of the first generation of stars (Pop III stars) are expected to host relatively few progenitor SNe, which could bearound ten, compared to systems with higher metallicity. In thiscase, the effects of stochastic sampling of an IMF may become
important.
The present sample of Boötes I stars are rather homogeneous in abundance ratios for elements lighter than zinc, which
does not strongly support the existence of any stochastic effect.
Nevertheless, the observed homogeneous [X/Fe] ratios do notcompletely rule out the stochastic enrichment scenario, since little scatter in [X/Fe] is expected, while producing large scatterin [Fe/H] in such scenarios as proposed by Frebel & Bromm(2012). This scenario thus implies that Boötes I could be a candidate for the first galaxies formed in the Universe.
Lee et al. (2013) have recently investigated the effect of thestochastic IMF sampling as a possible interpretation of the abundance difference and similarity between the UFDs and MW halostars. In their model, a system enriched by a smaller number of aprior generation of stars is expected to show a lower mean valueof [X/Fe] with larger spread, when the supernova yield of anelement X has a positive mass dependence (i.e., higher ejectedmass for a higher supernova progenitor mass). Therefore, if a
UFD has been enriched by fewer SNe (∼ 1 − 7 on average) thanthe progenitors of the MW halo (> 7) and if a Sr yield, which
is currently uncertain, has a strong positive mass dependence,
we would expect lower [Sr/Fe] for the UFDs than the MW haloprogenitor. They conclude that the stochastic sampling of IMF
for UFDs can simultaneously explain the observed difference in[Sr, Ba/Fe] ratios and the similarity in [Ti/Fe] ratios between the
UFDs and MW halo stars.
The result of the present work, a modest difference in[Mg/Fe] and a large difference in [Sr/Fe] from field halo stars
appear to both be reproduced with the above model if the Sryield depends more strongly on progenitor masses than the Mg
yields do. Thus, according to this model, the difference in number of prior-generation stars can explain the observed abundance
difference between the Boötes I and progenitors of the MW halo.
5.3.5. Individual yields

We finally discuss whether the abundances of the present sample
stars are consistent with yields of a single SN of Pop III stars.
Pop III stars may have been very massive. Very massive stars
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A&A proofs: manuscript no. boo16

with M> 140M are predicted to be exploded as pair-instability
SNe whose yields are expected to show a peculiar abundance
pattern, including very large Ca abundances. As can be seen,none of our sample stars show this peculiar pattern and seem to
be more consistent with yields of core-collapse SNe (e.g., Tominaga et al. 2007).
An intriguing issue is whether nucleosynthetic yields of SNeare compatible with the observed abundances of Sr and Ba in
Boötes I. Core-collapse SNe of massive (> 10M ) stars havebeen considered as one of the promising production sites ofheavy neutron-capture elements such as Ba (e.g., Thielemannet al. 2011). However, more recent studies report difficulty inreproducing the required physical condition for the r-process tooccur in the environment of the core-collapse SNe.
Qian & Wasserburg (2008) used the phenomenologicalmodel to study nucleosynthetic yields of neutron-capture elements in neutrino-driven winds, which are associated with SNe.
They pointed out that the low [Sr/Fe] ratios are not reproducedwith the predicted yields of Type II SNe with progenitor massesof neither 8 − 10M (H source) nor 12 − 25M (L sources).
Instead, they suggest that another source that produces a large
amount of Fe but little Sr is required to explain the abundance
patterns of stars with [Sr/Fe]<< −0.32. They speculate that anenergetic SN, a so-called hypernova, could be a possible astrophysical site that yields the low [Sr/Fe] ratios. In this scenario,
the observed Boötes I stars may have formed out of interstellar medium predominantly enriched by a single supernovae withlarge explosion energy.
A merging of neutron stars in a close binary system is proposed as an alternative site of the r-process. Recent calculationsof nucleosynthetic yields for the neutron-star merger successfully explain the abundance patterns of known r-process-richstars in the Galactic field halo (e.g., Wanajo & Janka 2012; Korobkin et al. 2012). It is interesting for future studies to investigate whether the predicted nucleosynthetic yields of neutron starmergers can explain the observed Sr/Fe and/or Ba/Fe abundancesin Boötes I.
The supernova yields depend on various unknown parameters, such as mass cut or explosion energies, and thus propertiesof the progenitor SN are not fully constraind at this moment. Nucleosynthetic yields of light and heavy neutron capture elements
for different SN progenitor masses at extremely low metallicities are required to examine the reason for the Sr deficiency in
Boötes I.

6. Conclusions
The chemical abundances of C, Na, α, iron-peak, and neutron
capture elements in six giant stars in Boötes I ultra-faint dwarf
galaxy were estimated based on the high-resolution spectra obtained with Subaru/HDS. Our main results are summarized asfollows:
1. The very low metallicity of Boo–094 ([Fe/H]= −3.4) previously reported in Feltzing et al. (2009) and Gilmore et al.
(2013) is confirmed. For the six sample stars, extremely
carbon-rich objects are not found.
2. For α and iron-peak elements, the scatter in the abundance
ratios is small over the metallicity range of −2.7 <[Fe/H]<−1.8.
3. The [Mg/Fe] and [Ca/Fe] ratios for these stars are slightly
lower than the field halo stars on average in the metallicity
range of −3.0 <[Fe/H]< −2.0. On the other hand, the abundance ratios for Na and Sc in the Boötes I stars are generally
similar to those of the field halo stars.
Article number, page 12 of 15page.15

4. The [Sr/Fe] abundances of the three Boötes I stars are lower
than the field halo values at similar metallicities. These
Boötes I stars deviate from the correlation between log (Ba)and log (Sr) seen in field MW halo stars and generally show
lower Sr/Ba ratios.
Highly inhomogeneous chemical abundances that are reported in fainter dSphs are not seen in our Boötes I sample.
Instead, our results suggest that the star formation in Boötes I
had lasted until Type Ia SNe started to contribute to the chemical enrichment, which has also been reported in other more luminous dSphs. Larger samples of chemical abundances and thekinematics of individual stars, as well as a determination of theorbital motion around the MW, are desirable to conclude whether
Boötes I is once more similar to brighter dSphs or is a surviving
example of first galaxies formed out of ejecta from Pop III stars.
Acknowledgements. We thank the referee for constructive comments and suggestions that have improved our paper. We are grateful to S. Feltzing for providing us with their Mg linelist. M.N.I. thanks K. Nomoto, L. Vargas, and A.
Frebel for helpful discussions on the abundance analysis and chemical evolutionof ultra-faint dwarf galaxies. M.N.I. acknowledges financial support from Grantin-Aid for JSPS (Japan Society for the Promotion of Science) fellow. W.A. was
supported by the JSPS Grants-in-Aid for Scientific Research (23224004).

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Article number, page 13 of 15page.15

A&A proofs: manuscript no. boo16
Table 4. Equivalent widths and abundances. Table 4 is published in its entirety in the electronic edition of the Astronomy & Astrophysics. A
portion is shown here for guidance regarding its form and context.

ObjectElem.

Ion

Boo–009
Boo–009
Boo–009

11
12
12

1
1
1

Wavelength(Å)
5895.92
4702.995528.40

χ
(eV)
0.00
4.35
4.35

log g f
−0.19
−0.44
−0.50

EW
(mÅ)
150.0
50.1
69.7

log A
(dex)
3.61
4.94
5.24

∆+100K−100K+0.3dex∆−0.3dex

Table 5. Elemental abundances and uncertainties

Name
Boo–009

Boo–094

Boo–117

Boo–121

Elem.
FeI
FeII
C(CH)
NaI
MgI
CaI
ScII
TiI
TiII
CrI
NiI
ZnI
SrII
BaII
FeI
FeII
C(CH)
NaI
MgI
CaI
ScII
TiI
TiII
CrI
NiI
ZnI
SrII
BaII
FeI
FeII
C(CH)
NaI
MgI
CaI
ScII
TiI
TiII
CrI
NiI
ZnI
SrII
BaII
FeI
FeII
C(CH)
NaI
MgI

log A
7.50
7.50
8.43
6.24
7.60
6.34
3.15
4.95
4.95
5.64
6.22
4.56
2.87
2.18
7.50
7.50
8.43
6.24
7.60
6.34
3.15
4.95
4.95
5.64
6.22
4.56
2.87
2.18
7.50
7.50
8.43
6.24
7.60
6.34
3.15
4.95
4.95
5.64
6.22
4.56
2.87
2.18
7.50
7.50
8.43
6.24
7.60

Article number, page 14 of 15page.15

log A
4.86
4.59< 5.50
3.61
5.09
3.88
0.39
2.56
2.31
2.25
3.60
< 2.20
...
−1.61
4.32
4.02
< 5.50
2.73
4.81
3.61
...
1.79
2.02
1.94
2.96
< 2.30
−2.47
−2.10
5.35
5.11
5.49
3.84
5.50
4.21
0.85
2.81
3.00
3.41
4.11
< 2.50
−0.40
−0.64
5.01
4.89
< 5.70
3.50
5.31

[X/H]
−2.64
−2.91
< −2.93
−2.63
−2.51
−2.46
−2.76
−2.39
−2.64
−3.39
−2.62
< −2.36
...
−3.79
−3.18
−3.48
< −2.93
−3.51
−2.79
−2.73
...
−3.16
−2.93
−3.70
−3.26
< −2.26
−5.34
−4.28
−2.15
−2.39
−2.94
−2.40
−2.10
−2.13
−2.30
−2.14
−1.95
−2.23
−2.11
< −2.06
−3.27
−2.82
−2.49
−2.61
< −2.73
−2.74
−2.29

[X/Fe]−2.64−2.91< −0.29
0.01
0.13
0.18
0.15
0.24
0.27
−0.75
0.02
< 0.28
...
−0.89
−3.18
−3.48
< 0.25
−0.32
0.39
0.46
...
0.02
0.55
−0.52
−0.08
< 0.92
−2.16
−0.80
−2.15
−2.39
−0.79
−0.25
0.04
0.01
0.08
0.01
0.44
−0.08
0.04
< 0.09
−1.12
−0.43
−2.49
−2.61
< −0.24
−0.25
0.20

N
38
2
1
1
2
2
2
2
2
1
2
1
0
1
22
2
1
1
3
2
0
1
1
1
1
1
1
3
32
3
1
2
3
6
3
3
1
1
1
1
1
2
34
4
1
2
1

σline
0.04
0.01
...
0.25
0.15
0.04
0.09
0.07
0.11
0.25
0.07
...
...
0.25
0.04
0.16
...
0.19
0.05
0.01
...
0.19
0.19
0.19
0.19
...
0.70
0.10
0.04
0.14
0.30
0.05
0.13
0.11
0.09
0.13
0.20
0.20
0.20
...
1.45
0.05
0.03
0.07
...
0.05
0.18

0.12
−0.12
0.01
0.00

...

0.01
−0.01
−0.06
0.06
−0.03
0.04
0.04
−0.04
0.01
−0.01
0.04
−0.04

−0.02
0.03
0.10
−0.09

...

−1

+0.3kms∆−0.3kms
−1
−0.04
0.06
−0.01
0.02

...

−0.03
0.02
0.00
−0.00
0.00
−0.00
−0.00
0.01
−0.00
0.00
−0.01
0.01
−0.00
0.00
0.00
−0.01

−0.10
0.10
0.01
−0.02
0.00
−0.01
0.00
−0.00
0.02
−0.02
−0.02
0.03

0.08
−0.08
0.12
−0.12

−0.01
0.01
−0.04
0.04
0.09
−0.08

−0.02
0.02
−0.02
0.03
−0.03
0.03

−0.01
0.00
−0.03
0.03
−0.04
0.05

−0.01
0.01
0.00
−0.00
0.01
−0.01

−0.07
0.08
0.01
−0.01
0.01
−0.01

0.01
−0.01
0.04
−0.04
0.00
−0.00
−0.00
0.00

0.00
−0.00
−0.00
−0.00
−0.00
0.00
0.00
−0.00

0.01
−0.02
0.02
−0.03
−0.00
0.00
0.00
−0.00

0.07
−0.08
0.13
−0.13
0.00
0.01
−0.13
0.13
0.02
−0.02
−0.03
0.03
−0.04
0.05
0.04
−0.04
0.01
−0.00
0.04
−0.04
−0.01
0.02
0.01
−0.01

−0.01
0.01
−0.03
0.03
0.10
−0.10
0.03
−0.03
−0.05
0.04
−0.00
0.00
0.01
−0.01
−0.00
0.00
0.00
−0.00
−0.01
0.01
0.01
−0.01
−0.01
0.01

0.01
−0.01
−0.10
0.13
−0.04
0.05
0.10
−0.13
−0.05
0.02
0.03
−0.05
0.05
−0.07
0.01
−0.01
0.04
−0.06
−0.02
0.03
0.08
−0.10
−0.06
0.07

0.07
−0.07
0.15
−0.16
−0.00
0.02

−0.00
0.00
−0.05
0.05
0.09
−0.08

−0.03
0.05
−0.08
0.11
−0.06
0.08

0.03
−0.05
−0.08
0.09

−0.03
0.02
0.00
−0.00

−0.08
0.05
0.02
−0.04

...

0.00
−0.00

...
...

...
...

...

...
...

...
...

...

...
...

...

...

...
...

...
...

...

...

0.01
−0.01

...
...

...

...

...
...

...
...

...

σtot
0.14
0.10
...
0.27
0.16
0.05
0.09
0.07
0.12
0.25
0.07
...
...
0.27
0.14
0.19
...
0.20
0.06
0.05
...
0.19
0.19
0.19
0.19
...
0.70
0.12
0.18
0.18
0.35
0.08
0.14
0.13
0.10
0.14
0.20
0.22
0.21
...
1.45
0.10
0.19
0.13
...
0.10
0.20

M. N. Ishigaki et al.: Chemical compositions of six metal-poor stars in the ultra-faint dwarf spheroidal galaxy Boötes I
Table 5. continued.

Name

Boo–127

Boo–911

HD216143

HD85773

Elem.
CaI
ScII
TiI
TiII
CrI
NiI
ZnI
SrII
BaII
FeI
FeII
C(CH)
NaI
MgI
CaI
ScII
TiI
TiII
CrI
NiI
ZnI
SrII
BaII
FeI
FeII
C(CH)
NaI
MgI
CaI
ScII
TiI
TiII
CrI
NiI
ZnI
SrII
BaII
FeI
FeII
C(CH)
NaI
MgI
CaI
ScII
TiI
TiII
CrI
NiI
ZnI
SrII
BaII
FeI
FeII
C(CH)
NaI
MgI
CaI

log A
6.34
3.15
4.95
4.95
5.64
6.22
4.56
2.87
2.18
7.50
7.50
8.43
6.24
7.60
6.34
3.15
4.95
4.95
5.64
6.22
4.56
2.87
2.18
7.50
7.50
8.43
6.24
7.60
6.34
3.15
4.95
4.95
5.64
6.22
4.56
2.87
2.18
7.50
7.50
8.43
6.24
7.60
6.34
3.15
4.95
4.95
5.64
6.22
4.56
2.87
2.18
7.50
7.50
8.43
6.24
7.60
6.34

log A
4.09
0.41
2.59
2.41
...
3.36
< 2.00
...
−1.06
5.58
5.58
5.74
4.14
5.79
4.61
1.02
3.12
3.23
3.61
4.21
< 2.70
−0.49
−0.61
5.34
5.32
5.50
3.80
5.79
4.17
...
2.60
2.79
2.83
3.90
< 2.20
...
−0.64
5.35
5.35
...
...
5.80
4.40
0.99
2.92
3.00
3.23
4.02
2.55
...
−0.09
5.11
5.17
...
3.99
5.45
4.11

[X/H]
−2.25
−2.73
−2.36
−2.54
...
−2.86
< −2.56
...
−3.24
−1.92
−1.92
−2.69
−2.10
−1.81
−1.73
−2.13
−1.83
−1.72
−2.03
−2.02
< −1.86
−3.36
−2.79
−2.16
−2.18
−2.93
−2.44
−1.81
−2.17
...
−2.35
−2.16
−2.81
−2.32
< −2.36
...
−2.82
−2.15
−2.15
...
...
−1.80
−1.94
−2.16
−2.04
−1.95
−2.41
−2.20
−2.01
...
−2.27
−2.39
−2.33
...
−2.25
−2.15
−2.23

[X/Fe]
0.24
−0.13
0.12
0.07
...
−0.37
< −0.07
...
−0.63
−1.92
−1.92
−0.77
−0.18
0.11
0.19
−0.21
0.09
0.20
−0.11
−0.10
< 0.06
−1.44
−0.87
−2.16
−2.18
−0.77
−0.28
0.35
−0.01
...
−0.19
0.02
−0.65
−0.16
< −0.20
...
−0.64
−2.15
−2.15
...
...
0.35
0.21
−0.01
0.11
0.20
−0.26
−0.05
0.14
...
−0.12
−2.39
−2.33
...
0.14
0.24
0.16

N
4
3
2
1
0
2
1
0
1
42
6
1
2
2
5
3
4
1
2
1
1
1
1
26
2
1
2
1
5
0
5
1
1
3
1
0
3
64
13
0
0
2
16
8
18
6
4
12
2
0
3
58
12
0
1
3
16

σline
0.06
0.11
0.19
0.18
...
0.02
...
...
0.18
0.03
0.07
0.30
0.14
0.01
0.09
0.05
0.09
0.22
0.04
0.22
...
1.04
0.22
0.02
0.06
0.30
0.05
0.11
0.09
...
0.08
0.11
0.11
0.17
...
...
0.12
0.01
0.01
...
...
0.07
0.03
0.03
0.03
0.04
0.03
0.02
0.01
...
0.06
0.01
0.01
...
0.11
0.08
0.02

∆+100K
−100K
−0.05
0.05
0.04
−0.04
0.01
−0.02
0.04
−0.04

...

+0.3dex∆−0.3dex
0.00
−0.01
0.00
−0.01
−0.00
0.01
0.00
−0.00

...

−1

+0.3kms∆−0.3kms
−1
0.02
−0.02
0.04
−0.05
0.03
−0.03
0.03
−0.04

...

−0.00
0.01

0.01
−0.01

0.03
−0.04

0.07
−0.08
0.15
−0.15
−0.01
0.02
−0.15
0.15
0.02
−0.02
−0.06
0.06
−0.04
0.04
0.03
−0.04
0.02
−0.02
0.03
−0.04
−0.01
0.00
0.02
−0.01

0.00
−0.01
−0.03
0.03
0.10
−0.10
0.03
−0.03
−0.06
0.06
−0.02
0.02
−0.01
0.00
0.00
−0.00
−0.00
0.01
−0.00
0.00
0.00
−0.00
−0.01
0.01

−0.01
0.02
−0.12
0.14
−0.10
0.13
0.12
−0.14
−0.02
−0.02
0.02
−0.04
0.02
−0.03
0.06
−0.08
−0.01
0.02
0.01
−0.01
0.05
−0.05
−0.07
0.08

0.06
−0.06
0.17
−0.18
−0.02
0.03
−0.17
0.18
0.04
−0.05
−0.07
0.08
−0.06
0.06

0.01
−0.01
−0.04
0.05
0.09
−0.08
0.04
−0.05
−0.03
0.03
−0.02
0.02
0.00
−0.01

0.07
−0.09
−0.12
0.15
−0.12
0.16
0.12
−0.15
−0.05
0.00
0.00
−0.02
0.04
−0.06

0.02
−0.03
0.04
−0.04
0.03
−0.04
−0.02
0.03

−0.00
0.00

0.05
−0.06
0.04
−0.06
−0.07
0.08
0.03
−0.04

0.08
−0.08
0.15
−0.16
−0.02
0.03

0.01
−0.01
−0.03
0.03
0.11
−0.10

−0.02
0.03
−0.07
0.09
−0.04
0.05

−0.05
0.05
−0.05
0.05
0.04
−0.04
0.02
−0.03
0.03
−0.04
0.01
−0.01
−0.04
0.05
−0.13
0.16

−0.04
0.04
−0.00
−0.00
−0.00
0.00
−0.01
0.00
−0.01
0.01
−0.01
0.01
0.02
−0.02
0.09
−0.09

−0.03
0.02
0.03
−0.04
0.01
−0.01
0.03
−0.04
−0.04
0.05
0.02
−0.02
0.05
−0.07
0.04
−0.05

0.07
−0.09
0.14
−0.16
−0.02
0.03

−0.01
0.01
−0.04
0.04
0.10
−0.10

−0.13
0.16
−0.06
0.07
−0.03
0.04

−0.08
0.09
−0.07
0.07
−0.05
0.05

0.02
−0.02
−0.02
0.01
0.00
−0.01

0.05
−0.07
−0.00
−0.01
0.02
−0.03

...
...

...
...

...

...
...

...
...

...

...

...
...

...
...

...

...

−0.02
0.02
0.01
−0.01

...
...

...
...

...

...

...
...

...
...

...

...
...

...
...

...

...

σtot
0.08
0.13
0.19
0.19
...
0.04
...
...
0.20
0.21
0.17
0.36
0.15
0.07
0.10
0.09
0.09
0.22
0.06
0.23
...
1.04
0.24
0.23
0.18
0.37
0.08
0.13
0.12
...
0.10
0.13
0.14
0.17
...
...
0.15
0.18
0.12
...
...
0.10
0.07
0.05
0.05
0.07
0.04
0.08
0.17
...
0.18
0.17
0.11
...
0.15
0.10
0.07

Article number, page 15 of 15page.15

A&A proofs: manuscript no. boo16
Table 5. continued.

Name

Elem.
ScII
TiI
TiII
CrI
NiI
ZnI
SrII
BaII

log A
3.15
4.95
4.95
5.64
6.22
4.56
2.87
2.18

Article number, page 16 of 15page.15

log A
0.74
2.73
2.89
3.03
3.85
2.67
...
−0.68

[X/H]
−2.41
−2.22
−2.06
−2.61
−2.37
−1.89
...
−2.86

[X/Fe]
−0.08
0.17
0.28
−0.22
0.02
0.50
...
−0.53

N
6
19
6
4
11
2
0
3

σline
0.03
0.02
0.04
0.02
0.02
0.00
...
0.07

∆+100K
−100K
0.04
−0.05
0.03
−0.05
0.04
−0.04
0.01
−0.02
−0.04
0.05
−0.13
0.17

...

0.09
−0.09

+0.3dex∆−0.3dex
−0.00
0.00
−0.01
0.00
−0.02
0.02
−0.01
0.01
0.02
−0.02
0.09
−0.09

...

−0.01
0.01

−1

+0.3kms∆−0.3kms
−1
0.01
−0.01
0.01
−0.01
−0.05
0.07
0.01
−0.01
0.04
−0.05
0.01
−0.01

...

−0.08
0.11

σtot
0.06
0.05
0.08
0.03
0.07
0.17
...
0.14