Issue |
A&A
Volume 519, September 2010
|
|
---|---|---|
Article Number | A14 | |
Number of page(s) | 7 | |
Section | Galactic structure, stellar clusters, and populations | |
DOI | https://doi.org/10.1051/0004-6361/201014135 | |
Published online | 07 September 2010 |
Light-element abundance variations in the Milky Way halo
S. L. Martell - E. K. Grebel
Astronomisches Rechen-Institut, Zentrum für Astronomie der Universität Heidelberg, 69120 Heidelberg, Germany
Received 26 January 2010 / Accepted 18 May 2010
Abstract
We present evidence for the contribution of high-mass globular
clusters to the stellar halo of the Galaxy. Using SDSS-II/SEGUE spectra
of over 1900 G- and K-type halo giants, we identify for the first time
a subset of stars with CN bandstrengths significantly larger, and CH
bandstrengths lower, than the majority of halo field stars, at fixed
temperature and metallicity. Since CN bandstrength inhomogeneity and
the usual attendant abundance variations are presently understood as a
result of star formation in globular clusters, we interpret this subset
of halo giants as a result of globular cluster dissolution into the
Galactic halo. We find that 2.5% of our sample is CN-strong, and can
infer based on recent models of globular cluster evolution that the
fraction of halo field stars initially formed within globular clusters
may be as large as 50%.
Key words: stars: abundances - Galaxy: halo - galaxies: formation
1 Introduction
Hierarchical structure formation is presently the dominant explanation for galaxy formation, based on observed fluctuations in the cosmic microwave background (Hinshaw et al. 2007) and sophisticated numerical simulations of their evolution to the present day (e.g., Diemand et al. 2007; Springel et al. 2008). In this picture, galaxies like the Milky Way are formed through the coalescence of multiple low-mass galaxies which develop within a much larger dark matter halo. The initial disagreement between the calculated mass function of dark matter subhaloes in the simulations and the observed mass function of nearby dwarf galaxies (e.g., Klypin et al. 1999; Moore et al. 1999) is being addressed from several directions at the same time, from complex semianalytic simulations that include star formation and feedback processes, and calculate the chemodynamical evolution of Milky Way-like galaxies (e.g., Johnston et al. 2008; Tumlinson 2010, and references therein) to searches for extremely low-mass galaxies in the Local Group (e.g., Zucker et al. 2006a,b; Belokurov et al. 2007).
The stellar halo of the Milky Way is thought to have been constructed mostly through the early (
Gyr
ago) accretion of low-mass protogalaxies. The halo exhibits
considerable substructure in density and in velocity (e.g. Bell
et al. 2008), and kinematically distinct streams presently observed in the halo (e.g., Majewski et al. 2003; Duffau
et al. 2006; Martínez-Delgado et al. 2007) are interpreted as remnants of more recent or ongoing merger activity. The ``ECHOS'' identified in Schlaufman et al. (2009)
are interpreted as older substructure that has lost some spatial
coherence with time. Studies of the abundance distributions and star
formation histories in nearby dwarf galaxies (e.g., Koch et al. 2007a,b, 2008a,b; Kirby et al. 2008; Aoki et al. 2009; Frebel et al. 2009, 2010)
have shown that there is a reasonable concordance between the
properties of the present-day Milky Way halo (as characterized by,
e.g., Schörck et al. 2009) and the dwarf
galaxies that would have been available as stellar contributors early
in Galactic history (e.g., Font et al. 2006; Carollo et al. 2007).
However, the ongoing dissolution of globular clusters such as Palomar 5 (e.g., Odenkirchen et al. 2001, 2002, 2003; Rockosi et al. 2002; Grillmair & Dionatos 2006) and NGC 5466 (Belokurov et al. 2006) implies that some fraction of halo stars are initially formed in globular clusters. This dissolution is driven by internal 2-body relaxation, stellar evolution processes and tidal interactions with the Galactic potential, and can cause significant mass loss over the lifetime of typical halo globular clusters (e.g., Gnedin & Ostriker 1997).
One particular model of globular cluster formation, described in D'Ercole et al. (2008),
posits that all old halo clusters surviving to the present day have
lost at least 90% of their initial mass. Early in the development of
the cluster, winds from AGB stars collected in the cluster center and
formed a second generation of low-mass stars. Type Ia supernovae then
caused the cluster to expand, and stars at large radii, mostly members
of the first generation, were lost. In the model, the end result,
Gyr later, is a
cluster with two stellar populations that differ slightly in age and abundance pattern.
This two-generation model was developed specifically to explain the
anomalous light-element abundance patterns observed in globular cluster
stars, specifically the presence in every old globular cluster in the
Milky Way of a subpopulation with typical Population II abundances, and
a
second subgroup with the same metallicity but enhanced N, Na, and Mg
along with depleted C, O,
and Al. This abundance bimodality has been studied extensively in
globular clusters (Langer
et al. 1992; Kraft 1994; and Gratton et al. 2004 all provide thorough reviews of the topic), and explanations for the second abundance subgroup have varied from
pollution (Cannon et al. 1998) to internal mixing (Langer 1985) to enrichment of star-forming gas by moderate- to high-mass stars (e.g., Cottrell & DaCosta 1981; Yong et al. 2008) or high-mass binaries (de Mink et al. 2010).
Surface pollution of already-formed stars would result in abundance
anomalies that are erased at first dredge-up, while current models of
deep mixing (e.g., Charbonnel & Zahn 2007)
indicate that it only begins to operate after first dredge-up. Both of
these processes have been effectively ruled out as explanations for
abundance bimodality by the presence of abundance variations at all
evolutionary phases in globular clusters (e.g., Briley et al. 2002; Harbeck et al. 2003).
In the primordial enrichment scenario, there is ongoing discussion over
the exact source of enriching material. In some models the source is
moderately high-mass
AGB stars (e.g., Parmentier et al. 1999), while Decressin et al. (2007) claim that rotating high-mass
stars are a better source for processed material because of their very short lifetimes and de Mink et al. (2010) prefer high-mass binaries because of their potentially strong mass loss and low wind velocities.
Stars with these light-element abundance anomalies are readily identified through strong UV/blue CN molecular absorption and relatively weak absorption in the CH G band, and are hereafter called ``CN-strong stars'', with the understanding that the full abundance pattern from C through Al necessarily follows the CN variation. They are not observed to exist in open clusters (e.g., Smith & Norris 1984; Jacobson et al. 2008; Martell & Smith 2009) or the halo field (Gratton et al. 2004). This feedback process apparently only occurs in the high-density environment of globular clusters, and as a result the characteristic sawtooth abundance pattern from carbon through aluminium can be used as a marker of globular cluster origin.
Given the contributions globular clusters are presently making to the halo field, and the significant mass loss predicted theoretically over the lifetime of the Galactic globular cluster system (e.g., Baumgardt et al. 2008), it is intriguing that no CN-strong stars have to date been observed in the halo. We interpret this as a qualitative sign that the contributions to the halo field of globular clusters as we know them today are relatively minor, as is also suggested in Yong et al. (2008).
To test this interpretation, we searched for CN-strong halo giants in
the Sloan Extension for Galactic Understanding and Exploration (SEGUE)
survey (Yanny et al. 2009). The SEGUE survey is a spectroscopic extension of imaging taken during the Sloan Digital Sky Survey (York et al. 2000),
with targets selected to address questions of halo substructure and
Galactic formation history. Data were taken from 2005 August through
2008 July using a 640-fiber multiobject spectrograph and the same
telescope at Apache Point Observatory that was used for SDSS imaging.
The first portion of the SEGUE data was made publically available in
2008 as part of SDSS Data Release 7 (DR7), including roughly
240 000 spectra in 200 ``pencil beam'' lines of sight containing
stars chosen for specific purposes (i.e., fields in the Sagittarius
stream, M dwarfs to study extremely local kinematic substructure, G and
K giants to study the distant halo). In addition to flux-calibrated
spectra, DR7 also offers the products of the ``SEGUE Stellar Parameters
Pipeline'' (SSPP), which include derived stellar parameters like
effective temperature, [Fe/H] metallicity, and radial velocity,
determined
automatically through template matching,
minimization, cross-correlation or grid-matching methods, as appropriate. Lee et al. (2008a,b) and Allende Prieto et al. (2008) give thorough explanations of the SSPP pipeline and process.
2 The data set
Data were obtained from SDSS DR7, through the online Catalog Archive
Server. We selected all SEGUE plates, and from those all stars with
[Fe/H] ,
,
,
and mean signal-to-noise per pixel (SNR) larger than 20 were chosen.
Additionally, the errors on various derived parameters were required to
be small:
,
with at least three independent [Fe/H] determinations, and reduced
of the best-fit template spectrum less than 2.0 both in the region of the Ca II H and K lines and in the CH G band.
This intentionally generous set of selection criteria then had to be
further sub-selected to isolate halo giants. To accomplish this
sub-selection, we lowered the limit on
to 3.0,
then divided the initial data set into 0.2-dex-wide bins in [Fe/H], and
removed AGB, main-sequence and turnoff stars from the H-R diagram of
each of those subsets by rejecting all points more than
in
from the fiducial sequence. We also made the SNR requirement more
stringent, requiring that the mean SNR per pixel in the wavelength
range
be larger than 15. This left 5066 halo giants, out of the original 22 784 stars. In Fig. 1,
the small points represent those halo giants, while density contours of
stars rejected based on CMD position are shown as solid lines.
![]() |
Figure 1:
Hertzsprung-Russell diagram for the initial SEGUE data set, split into
two subgroups. 5066 G and K giants are shown as small points,
while density contours for the remaining 17 718 stars rejected
based on (
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To facilitate later analysis, we convert the dereddened apparent
magnitudes given in the SEGUE data to absolute
magnitudes through a simple photometric parallax calculation. We created a grid of 10 Gyr Padova isochrones (Marigo et al. 2008)
with metallicities ranging from [Fe/H] =-1.0 to [Fe/H] =-1.8
(the lowest metallicity available) at a spacing of 0.1 dex, and
interpolate between those to match the metallicity of each individual
star. The absolute
magnitude corresponding to the observed
color of the star, on the interpolated isochrone, is then assigned as the true absolute
magnitude.
Errors in
were calculated by Monte Carlo sampling of uncorrelated errors in SSPP metallicity,
magnitude and
magnitude, drawn randomly from a Gaussian with a width equal to the
reported errors on those quantities. This addition of error was done 104 times for each star, and we take the standard deviation in those 104 determinations of
as the error on
.
This error has a typical value of 0.3 mag, with the largest values (
mag) on the
of stars with the largest errors in apparent
and
.
The age of the isochrones used had minimal effects on the derived
values, with a change of only
mag for a shift of
Gyr.
The limited metallicity range of the isochrones is sufficient for the
purposes of the analysis in Sect. 3, but lower-metallicity
isochrones would allow us to study the distance distribution of the
roughly 1/3 of our final dataset at lower metallicity, to compare the
spatial distribution of our final data set to the inner and outer halos
identified in Carollo et al. (2007).
We measured S(3839) (Norris et al. 1981), a bandstrength index for the CN band at
,
for all halo giant spectra. S(3839)
measures the magnitude difference between the integrated flux in the CN
feature and the integrated flux in a nearby continuum band, with more
absorption in the feature resulting in larger bandstrength. As can be
seen in Fig. 2, there is a strong concentration at low S(3839)
in our data set. This is to be expected, since the halo is primarily
composed of CN-weak stars with typical Pop. II abundances. There
is also a clear trend with temperature, in the sense that cooler stars
have larger CN bandstrengths. The cooling of stars as they ascend the
giant branch reddens the spectra and permits more CN molecule
formation, both of which increase the flux difference between the
feature and continuum bands of S(3839). There are, however, interesting outliers in Fig. 2: many of the stars with dramatically large S(3839) are carbon stars (shown as open triangles), with correspondingly large CH and C2bandstrengths. Figure 3
shows four sample spectra from the halo giant data set: the uppermost
spectrum (of SDSS J035123.90+092451.3) is a low-metallicity carbon
star, the next (of SDSS J115934.87+002748.0) is a possible CEMP star
(Carbon-Enhanced Metal-Poor, having [Fe/H]
and [C/Fe]
,
and described in Lucatello et al. 2006
and references therein), with strong CH absorption redward of the G
band and a low metallicity, the next spectrum (of SDSS
J064411.96+275351.6) is not a carbon star, but is CN-strong, and the
lowest spectrum (of SDSS J145301.24-001954.1) is a typical CN-weak halo
star.
![]() |
Figure 2: CN bandstrength index S(3839) versus dereddened (g-r) color for the G and K giants indicated in Fig. 1. As expected, the stars are mainly CN-weak, and there is a slight temperature dependence. As mentioned in the text, the stars with dramatically high S(3839) are mainly CEMP and carbon stars (shown as open triangles). |
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![]() |
Figure 3:
Sample spectra from the four main types of star comprising our G and K
giant data set: the topmost is from a carbon star, with clear CH and C2
bands, and the next is a possible CEMP star, with a metallicity of
[Fe/H] = -2.2 and strong CH absorption. The next spectrum down is
from a CN-strong star, with stronger absorption in the
|
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The broad molecular features in the example carbon-star spectrum are quite clear. We use the strength of the CH feature around
and the Swan (1, 0) C2 band at
to identify carbon and CEMP stars and remove them from the final data set. Specifically, we measure these indices:

and consider all stars with [Fe/H]


![$-1.8 \le [{\rm Fe/H}]\le -1.4$](/articles/aa/full_html/2010/11/aa14135-10/img43.png)


![$[{\rm Fe/H}]\ge -1.4$](/articles/aa/full_html/2010/11/aa14135-10/img46.png)



![$[{\rm Fe/H}]=-1.8$](/articles/aa/full_html/2010/11/aa14135-10/img49.png)
![$-1.8 \le [{\rm Fe/H}]\le -1.4$](/articles/aa/full_html/2010/11/aa14135-10/img43.png)
![$[{\rm Fe/H}]=-1.4$](/articles/aa/full_html/2010/11/aa14135-10/img50.png)
3 Distribution of CN and CH bandstrengths
Since the purpose of this study is identifying stars in the halo with
the CN-strong chemical signature associated with globular clusters, and
our data source is the moderate-resolution spectroscopy provided by
SEGUE, our methods follow closely the techniques used in low-resolution
spectroscopic studies of CN bandstrength behavior in globular clusters.
For example, the spectral feature used is the same
CN band, and the index used to measure bandstrength is the well-known S(3839).
One first check of this approach is comparing CN behavior in
well-studied globular clusters to halo stars with similar metallicities
observed as part of SEGUE. Figure 4
shows two views of the CN bandstrength data for M 3, measured from
SEGUE spectra of M 3 giants (membership information from J.
Smolinski, private communication). The left panel shows S(3839) versus absolute
magnitude, and the typical globular cluster pattern is readily visible: the S(3839)
distribution separates into two parallel loci, rising with increasing
luminosity. CN-strong stars (shown as open circles) are clearly
distinct from CN-weak stars (filled circles) at fixed luminosity.
![]() |
Figure 4:
S(3839) versus
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The right panel of Fig. 4 converts the raw S(3839)
data into a generalized histogram. In order to remove the
temperature-related trend visible in the left panel, we fit a line to
the CN-weak locus and measure the quantity
,
the vertical difference in S(3839)
between every point and the baseline. We then draw a generalized
histogram by representing each point as a Gaussian centered at
,
with a FWHM equal to the error on
,
and summing the individual Gaussians together. In Fig. 4, we base the Gaussian widths in the generalized histogram on the actual measurement errors on S(3839), as determined by Monte Carlo sampling of the spectral error vectors reported by SEGUE. Typical values of
are 0.014 mag, and we amplify that error by a factor of four in
the construction of the generalized histogram. This allows for errors
not accounted for in our Monte Carlo noise sampling, particularly
errors in flux calibration and the SSPP-derived parameters. We
calculate generalized histograms for the CN-strong and CN-weak stars
independently, and normalize both curves to the peak of the CN-weak
curve, so that the relative heights represent the relative numbers of
stars in the two groups.
Figure 5 shows analogous data for field stars in the final data set with [Fe/H] within dex of M 3, and similar structure can be seen to Fig. 4, with interesting differences in scale. As in Fig. 4, we fit a baseline to the CN-weak stars in the
plane (left panel), measure the quantity
,
and convert that into a generalized histogram in the right panel. Errors in
were calculated as in Fig. 4.
In this instance, we divide the CN-weak from CN-strong stars by
shifting the baseline vertically until it encounters a significant gap.
As in Fig. 4, CN-weak stars are
shown as filled circles, and CN-strong stars as open circles. Although
the two peaks in the right panel have roughly the same separation as in
Fig. 4, the CN-weak group is a
much more dominant component of the overall population. We have
magnified the CN-strong curve in the right panel by a factor of 30
(noted in the upper right corner of the panel) so that it is more
easily comparable to the CN-weak curve.
![]() |
Figure 5:
S(3839) versus
|
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It must be noted that the slope of the baseline in Fig. 5 is different from that in Fig. 4.
This is unexpected, since (at fixed metallicity) the same processes
ought to be moderating the progressive increase in CN bandstrength with
increasing giant-branch luminosity: a declining temperature that both
shifts flux from the science band of S(3839) into the continuum band and permits more CN molecule formation, along with the ``canonical extra mixing'' decribed
in Denissenkov & VandenBerg (2003)
that progressively depletes carbon and enhances nitrogen abundance in
the photosphere. We attribute the difference in baseline slopes to the
different mass distributions of the two samples. Specifically, the
M 3 giants are an old, single-age population, with masses all
around
,
while the field giants have the possibility of being considerably
younger, and therefore more massive. A younger and more massive red
giant, at fixed luminosity, will undergo more rapid evolution, leaving
less time for deep mixing and surface abundance changes. Indeed, as is
discussed in Gilroy (1989), stars with masses greater than
evolve along the giant branch too quickly to even begin deep mixing. As
a result, younger giants in the field will have weaker CN bands at a
fixed
,
with the effect of pulling down the overall baseline slope.
To extend this type of analysis to the full final data set, we must
carefully isolate [C/Fe] and [N/Fe] variations from the underlying
metallicity, which strongly affects the appearance of double-metallic
absorption features like
CN. To that end, we divide the 1958 stars in the final data set into
0.1-dex-wide bins in [Fe/H] and search the bins independently for a
bimodal distribution of CN bandstrength at fixed luminosity. Since CN
bandstrength variations become very small at low metallicity, even for
significant variations in [C/Fe] and [N/Fe] (e.g., Martell et al. 2008b; Briley et al. 1993), we limit our sample to the relatively metal-rich end of the halo. Figure 6 shows the raw S(3839) vs.
distribution for field stars in each of the metallicity bins between
and
.
The maximum metallicity of the bin is given in the upper left corner of each panel. As in the left panel of Fig. 4, the dashed line in each panel is the baseline against which
is measured.
![]() |
Figure 6:
Raw S(3839) vs.
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There are several features of note in this figure: the slope of the
baseline flattens monotonically with dropping metallicity, from -0.11
in the
bin, to -0.04 in the
bin. It is known that deep mixing is more efficient at low metallicity: Martell et al. (2008c) find that d[C/Fe]/dMV is twice as large at
as at -1.0, and Sweigart & Mengel (1979)
predict less-compressed hydrogen-burning shells, allowing for more
penetration by meridional circulation currents, at lower [Fe/H]. In
addition, the gap between the CN-strong and CN-normal groups shrinks as
metallicity declines, an effect that occurs even without reducing the
size of variations in [C/Fe] and [N/Fe].
In order to identify stars with globular cluster-like carbon and nitrogen abundances in Fig. 6,
we look for the classic CN-CH anticorrelation, used in globular cluster
studies as a clear signal of carbon depletion and nitrogen enrichment.
Figure 7 shows
versus the CH bandstrength index
(Martell et al. 2008a) for the subset of the final data set with
.
The stars are further subdivided by absolute
magnitude, since CH bandstrength is a sensitive function of both
temperature and carbon abundance, and both decrease with rising
luminosity. Stars with relatively large
and relatively weak
in each luminosity sub-bin are shown as open circles.
![]() |
Figure 7:
CN versus CH bandstrength for the
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We repeated this selection process, dividing the metallicity bins into
luminosity sub-bins, for the other seven metallicity bins, and
Fig. 8 shows
versus
for each of the metallicity bins, with all luminosity sub-bins
collapsed together. Stars with strong CN (and weak CH, in the
lower-metallicity bins), relative to the majority of halo field stars,
are shown as open circles in each panel. Altogether, we identify
49 stars (also shown as open circles in Fig. 8) as relatively CN-strong and CH-weak. Our CH bandstrength index
is calibrated based on bright red giants in the low-metallicity
globular cluster M 53. It is therefore less responsive to
variations in carbon abundance in high-metallicity stars and in fainter
giants than it was designed for. However, it was the most responsive
over the full parameter range of our data, of the nine G-band indices
we considered. As a result, we are less stringent in our CH
bandstrength selection in the higher-metallicity bins than in the
lower-metallicity bins.
![]() |
Figure 8: CN versus CH bandstrength for all eight metallicity bins, selected from luminosity sub-bins as in Fig. 7. Candidate CN-strong stars are shown as open circles. |
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![]() |
Figure 9:
Generalized histograms of
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The eight panels of Fig. 9 correspond to the panels in Fig. 6, and show generalized histograms of
for the CN-normal and CN-strong stars in each metallicity bin, using the CN bandstrength classifications made in Fig. 8. As in Fig. 5,
the CN-strong curves needed to be amplified to be easily visible
next to the CN-weak curves; the multiplicative factor is given in the
upper right corner of each panel, and the maximum metallicity in each
bin is given in the upper left corner of each panel. These generalized
histograms are qualitatively similar to Fig. 4,
although with a clearly different CN-strong/CN-weak ratio: there is a
separation between the two peaks of roughly 0.2 mag in
,
which shrinks as the overall metallicity drops.
The presence of stars in the halo field with relatively strong
CN absorption and relatively weak absorption in the
CH G band, at fixed metallicity and luminosity, is a strong indication
that globular clusters have contributed stars to the halo field. Given
current models for the origin of light-element abundance variations in
globular clusters, it does not seem possible that these stars formed in
the halo with these atypical [C/Fe] and [N/Fe] abundances. An
investigation of the abundances of O, Na, Mg and Al, which are also
known to vary in CN-strong globular cluster stars (see, e.g., Kraft 1994),
would allow a more strict test, and possibly a more firm confirmation,
of our claim that the 49 halo field stars we identify here as CN-strong
originated in globular clusters.
For completeness, we also mention the possibility that these stars did
not originate within globular clusters, but rather that their unusual
abundances are the result of mass transfer from an AGB companion in a
binary system. However, as is demonstrated in Lucatello et al. (2005), which discusses binarity in CEMP stars, the fraction of field stars expected to be in a binary system with a companion of
and orbital parameters that permit evolution of the companion up to the
AGB phase, and then mass transfer but not coalescence, is quite small.
There is no high-precision radial velocity monitoring program in
progress for these stars; we predict that such a program would be
unlikely to find binary companions.
4 Discussion
Although light-element abundance variations have not been observed
before in the halo, our identification of these CN-strong halo field
stars is not wholly unexpected. There are several well-understood
mechanisms for globular cluster mass loss, and theoretical studies of
globular
cluster formation and evolution (e.g., D'Ercole et al. 2008; Baumgardt et al. 2008)
predict significant mass loss in individual clusters as well as a
dramatic
reshaping of the cluster mass function with time. The data set we
analyzed, selected from the SEGUE survey, is not representative of the
full halo, in mass, evolutionary phase, or metallicity. Distances to
the candidate CN-strong stars range from 4 to nearly 40 kpc, with
found within 20 kpc of the Sun. However, our result for red giants
is generalizable to all halo stars, since abundance bimodality is
observed to exist at all masses and evolutionary phases in globular
clusters. More fundamentally, all of the cuts we made in selecting the
final data set are blind to CN and CH bandstrengths and light-element
abundances, and all nitrogen-enhanced giants with metallicities above
ought to show clearly strong CN bands. Since approximately
of our halo red giants exhibit strong CN bands and weak CH bands, we
expect that the same fraction of the entire halo will contain the same
abundance enhancements and depletions. This prediction can be confirmed
by observations of dwarfs in the halo field, because main sequence
stars in globular clusters show the same abundance division as giants
(e.g., Briley et al. 2002).
In order to convert
,
the present-day fraction of CN-strong halo stars, into
,
the fraction of globular cluster-originating stars in the halo field,
we must consider what fraction CN-strong stars comprise of the stars
originally formed in
globular clusters. In the two-generation scenario of D'Ercole et al. (2008), roughly
of stars originally formed in a globular cluster, consisting entirely
of first-generation stars (with halo-like chemistry), are lost between
the epoch of cluster formation and the present day. This means that
,
the fraction of stars that remain as members of the globular cluster they were formed in, is around 0.1. Since
,
the fraction of present-day globular clusters stars that are CN-strong, is around 0.5 (e.g., Kraft
1994), we can calculate that
.
Since
in the present study, this suggests that
,
and that a remarkable
of the halo field originally formed in the massive star clusters that
were progenitors of the present-day globular cluster population, with a
further unknown contribution of CN-weak stars made
by globular clusters that did not survive to the present day and were
not massive enough to self-enrich.
While some numerical studies of galaxy formation (e.g., Boley et al. 2009) have suggested that the halo could be constructed entirely from disrupted globular clusters, there is not presently a strong consensus on the role of globular clusters in cosmological-scale galaxy formation. Precise numerical study of the dynamical evolution of globular clusters is very complicated: the number of particles is large enough, and the relevant timescales short enough, that highly accurate simulations are very time-consuming. However, the development of semianalytic prescriptions for the mass evolution of globular clusters would allow single-halo-scale simulations like those of Johnston et al. (2008) to include them as a source of halo stars, and to predict what fraction of the halo field ought to originate in globular clusters.
AcknowledgementsS.L.M. wishes to thank Graeme Smith and Tim Beers for helpful conversations about this project. Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the US Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington.
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All Figures
![]() |
Figure 1:
Hertzsprung-Russell diagram for the initial SEGUE data set, split into
two subgroups. 5066 G and K giants are shown as small points,
while density contours for the remaining 17 718 stars rejected
based on (
|
Open with DEXTER | |
In the text |
![]() |
Figure 2: CN bandstrength index S(3839) versus dereddened (g-r) color for the G and K giants indicated in Fig. 1. As expected, the stars are mainly CN-weak, and there is a slight temperature dependence. As mentioned in the text, the stars with dramatically high S(3839) are mainly CEMP and carbon stars (shown as open triangles). |
Open with DEXTER | |
In the text |
![]() |
Figure 3:
Sample spectra from the four main types of star comprising our G and K
giant data set: the topmost is from a carbon star, with clear CH and C2
bands, and the next is a possible CEMP star, with a metallicity of
[Fe/H] = -2.2 and strong CH absorption. The next spectrum down is
from a CN-strong star, with stronger absorption in the
|
Open with DEXTER | |
In the text |
![]() |
Figure 4:
S(3839) versus
|
Open with DEXTER | |
In the text |
![]() |
Figure 5:
S(3839) versus
|
Open with DEXTER | |
In the text |
![]() |
Figure 6:
Raw S(3839) vs.
|
Open with DEXTER | |
In the text |
![]() |
Figure 7:
CN versus CH bandstrength for the
|
Open with DEXTER | |
In the text |
![]() |
Figure 8: CN versus CH bandstrength for all eight metallicity bins, selected from luminosity sub-bins as in Fig. 7. Candidate CN-strong stars are shown as open circles. |
Open with DEXTER | |
In the text |
![]() |
Figure 9:
Generalized histograms of
|
Open with DEXTER | |
In the text |
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