![]() | Annu. Rev. Astron. Astrophys. 2009. 47:
371-425 Copyright © 2009 by Annual Reviews. All rights reserved |
The detailed chemical abundance patterns in individual stars of a stellar population provide a fossil record of chemical enrichment over different timescales. As generations of stars form and evolve, stars of various masses contribute different elements to the system, on timescales directly linked to their mass. Of course, the information encoded in these abundance patterns is always integrated over the lifetime of the system at the time the stars studied were born. Using a range of stars as tracers provides snapshots of the chemical enrichment stage of the gas in the system throughout the SFH of the galaxy. This approach also assumes that the chemical composition at the stellar surface is unaffected by any connection between interior layers of the star, where material is freshly synthesised, and the photosphere. This assumption is generally true for main-sequence stars, but evolved stars (giants or super-giants) will have experienced mixing episodes that modified the surface composition of the elements involved in hydrogen burning through the CNO cycle, i.e. carbon, nitrogen and possibly also oxygen.
These studies require precise measurements of elemental abundances in individual stars and this can only be done with high-resolution and reasonably high signal-to-noise spectra. It is only very recently that this has become possible beyond our Galaxy. It is efficient high-resolution spectrographs on 8-10m telescopes that have made it possible to obtain high resolution (R > 40000) spectra of RGB stars in nearby dSphs and O, B and A super-giants in more distant dIs. These stars typically have magnitudes in the range V = 17-19. Before the VLT and Keck, the chemical composition of extra-galactic stars could only be measured in super-giants in the nearby Magellanic Clouds (e.g., Wolf 1973, Hill, Andrievsky & Spite 1995, Hill, Barbuy & Spite 1997, Venn 1999), yielding present day (at most a few 107yr ago) measurements of chemical composition. Looking exclusively at young objects however makes it virtually impossible to uniquely disentangle how this enrichment built up over time.
4.1. Dwarf Spheroidal Galaxies
The first studies of detailed chemical abundances in dSph galaxies are those of Shetrone, Bolte & Stetson (1998), Shetrone, Côté & Sargent (2001, 17 stars in Draco, Ursa Min & Sextans) using Keck-HIRES and Bonifacio et al. (2000, 2 stars in Sgr) using VLT-UVES. These early works were shortly followed by similar studies slowly increasing in size (Shetrone et al. 2003, Bonifacio et al. 2004, Sadakane et al. 2004, Geisler et al. 2005, McWilliam & Smecker-Hane 2005). The total number of stars probed in individual studies remained very low (typically only 3 to 6 in any one galaxy except for Sgr). This was because the stars had to be observed one at a time, and for the most distant dSphs this required exposure times of up to 5 hours per star. Nevertheless from these small samples it was already clear that dSph galaxies follow unique chemical evolution paths, which are distinct from that of any of the MW components (e.g., Shetrone, Côté & Sargent 2001, Shetrone et al. 2003, Tolstoy et al. 2003, Venn et al. 2004a).
Most recently, high-resolution spectrographs with high multiplex capabilities have resulted in large samples (> 80 stars) of high resolution spectra of individual stars to determine abundances in a relatively short time. The FLAMES multi-fiber facility on VLT (Pasquini et al. 2002) has so far been the most productive in this domain. There are a number of FLAMES high resolution spectroscopy studies in preparation, but some results are already published for Sgr and its stream (Monaco et al. 2005, Sbordone et al. 2007, Monaco et al. 2007, 39 stars), Fnx (Letarte 2007, 81 stars), Carina (Koch et al. 2008a, 18 stars) and Scl (Hill et al. in preparation, 89 stars).
These new extensive studies not only provide abundances with better statistics, but they also allow statistical studies over the total metallicity range in each galaxy. This allows for an almost complete picture of their chemical evolution over time, with abundance trends as a function of metallicity for each system. Only the most metal-poor regime in these systems is perhaps still somewhat under-represented in these samples, although this is in part because they are rare (Helmi et al. 2006), and in part because these large samples of abundances have been chosen in the inner parts of the galaxies, where younger and/or more metal-rich populations tend to dominate (Tolstoy et al. 2004, Battaglia et al. 2006). New studies to fill in this lack of measured abundances in low metallicity stars are in preparation (e.g., Aoki et al. 2009). In the following, we will consider groups of elements that give particular insights into dwarf galaxy evolution.
4.1.1 ALPHA ELEMENTS
The -elements abundances
that can easily be measured in RGB
spectra includes O, Mg, Si, Ca and Ti. Although the
-elements
have often been considered as an homogeneous group, and their
abundances are sometimes averaged to produce a single
[
/Fe]
ratio, their individual nucleosynthetic origin is not always exactly
the same. For example, O and Mg are produced during the hydrostatic He
burning in massive stars, and their yields are not expected to be
affected by the SNII explosion conditions. On the other hand Si, Ca and
Ti are mostly produced during the SNII explosion. This distinction is
also seen in the observations (e.g.,
Fulbright,
McWilliam & Rich 2007),
where Si, Ca and Ti
usually track one another, but O and Mg often show different trends
with [Fe/H]. It is therefore generally advisable to treat the three
-elements which are well
probed in dwarf galaxies separately.
Fig. 11 shows a compilation of Mg and Ca
abundances of individual stars in those dSphs with more than 15
measurements.
![]() |
Figure 11. Alpha-elements (Mg and Ca) in four nearby dwarf spheroidal galaxies: Sgr (red: Sbordone et al. 2007, Monaco et al. 2005, McWilliam & Smecker-Hane 2005), Fnx (blue: Letarte 2007, Shetrone et al. 2003), Scl (green: Hill et al. in prep; Shetrone et al. 2003, Geisler et al. 2005) and Carina (magenta: Koch et al. 2008a, Shetrone et al. 2003). Open symbols refer to single-slit spectroscopy measurements, while filled circles refer to multi-object spectroscopy. The small black symbols are a compilation of the MW disk and halo star abundances, from Venn et al. (2004a). |
The apparent paucity of
-elements (relative to
iron) in dSph galaxies compared to the MW disk or halo was first noted by
Shetrone, Bolte
& Stetson (1998),
Shetrone,
Côté & Sargent (2001),
Shetrone et
al. (2003),
Tolstoy et
al. (2003),
Venn et al. (2004a)
from small samples. Fig. 11 shows this
convincingly over most of
the metallicity range in each system. However, it also appears that
each of these dSphs starts, at low [Fe/H], with
[
/Fe] ratios
similar to those in the MW halo at low metallicities. These ratios
in the dSphs then evolve down to lower values
than is seen in the MW at the same metallicities.
The ratio of -elements
to iron, [
/Fe], is commonly
used to trace the star-formation timescale in a system, because it is
sensitive to the ratio of SNII (massive stars) to SNIa (intermediate
mass binary systems with mass transfer) that have occurred in the
past. SNIa have a longer time scale than SNII and as soon as they
start to contribute they dominate the iron enrichment
and [
/Fe] inevitably
decreases. After that, no SFH can ever again result in enhanced
[
/Fe], unless coupled
with galactic winds removing only the SNIa ejecta and not that of SNII.
This is seen as a "knee" in a plot of [Fe/H] vs.
[
/Fe], see
Fig. 11. The knee position indicates the
metal-enrichment achieved by a system at the time SNIa start to
contribute to the chemical evolution (e.g.,
Matteucci &
Brocato 1990,
Matteucci 2003).
This is between 108 and 109yrs after the first star
formation episode. A galaxy that efficiently produces and
retains metals over this time frame will reach a higher metallicity
by the time SNIa start to contribute than a galaxy which either loses
significant metals in a galactic wind, or simply does not have a very
high SFR. The position of this knee is expected to be different for
different dSphs because of the wide variety of SFHs. In the data
there are already strong hints that not all dSphs have a knee at the
same position.
At present the available data only cover the knee with sufficient
statistics to quantify the position in the Scl dSph, a system
which stopped forming stars 10 Gyr ago, and the knee occurs at
[Fe/H] -1.8. This is
the same break-point as the two
kinematically distinct populations in this galaxy
(Tolstoy et
al. 2004,
Battaglia
2007),
see Fig. 9. This means that the
metal-poor population has formed before any SNIa enrichment took
place, which means on a timescale shorter than 1 Gyr.
In other dSphs the knee is not well defined due to a lack of data, but
limits can be established. The Sgr dSph has enhanced
[/Fe] up to [Fe/H]
-1.0, which is
significantly more
metal-rich than the position of the knee in the Scl dSph. This is
consistent with what we know of the SFH of Sgr, which has
steadily formed stars over a period of 8-10 Gyrs, and only stopped
forming stars about 2-3 Gyr ago (e.g.,
Dolphin 2002).
The Carina dSph has had an unusually complex SFH, with
at least three separate bursts of star formation
(Hurley-Keller,
Mateo & Nemec 1998),
see Fig. 4. The
abundance measurements in Carina are presently too scarce to have any
hope to confidently detect these episodes in the chemical enrichment
pattern (e.g.,
Tolstoy et
al. 2003).
It appears to possesses
[
/Fe] poor stars between
[Fe/H] = -1.7 and -2.0, which
suggests that the knee occurs at lower [Fe/H] than in Scl. It seems that Carina has had the least amount of chemical evolution
before the onset of SNIa of all galaxies in
Fig. 11. In
the Fnx dSph, another galaxy with a complex SFH, the
sample does
not include a sufficient number of metal-poor stars to determine even
an approximate position of the knee. There are abundances for only five
stars below [Fe/H] = -1.2, and only one below [Fe/H] = -1.5. The
knee is constrained to be below [Fe/H] < -1.5. From this (small)
sample of dSph galaxies, it appears that the position of the knee
correlates with the total luminosity of the galaxy, and the mean
metallicity of the galaxy. Which suggests that the presently most
luminous galaxies are those that must have formed more stars at the
earliest times and/or retained metals more efficiently than the less
luminous systems.
The abundance ratios observed in all dSphs for stars on the metal-poor
side of the knee, tend to be indistinguishable from those in the
MW halo. From this small sample it seems that the first
billion years of chemical enrichment gave rise to similar enrichment
patterns in small dwarf galaxies and in the MW halo. Because
the [/Fe] at early times
is sensitive to the IMF of the massive stars, if
[
/Fe] in metal-poor
stars in dSphs and in
the MW halo (or even the bulge) are similar, then there is no
need to resort to IMF variations between these systems. Poor IMF
sampling has been invoked as a possible cause of lowering
[
/Fe] in dwarfs
(Tolstoy et
al. 2003,
Carigi &
Hernandez 2008),
but these new large samples suggest that this explanation may no longer
be necessary, at least in systems as luminous as Sculptor, Fornax or
Sagittarius.
On the other hand, there is now a hint that the slightly less
luminous Sextans (MV = -9.5) could display a
scatter in the
/Fe ratios at the
lowest metallicities, including
[
/Fe] close to solar
(Aoki et al. 2009).
Such a scatter is so far
observed only in this purely old system, and suggests a
very inhomogeneous metal-enrichment in this system that presumably
never retained much of the metal it produced. The true extent of this
scatter in
Sextans remains to be investigated, and extension to
other similar
systems is needed before general conclusions can be reached on the
mechanisms leading to the chemical homogeneity -or not- of dwarf galaxies.
At later times, in those stars which formed ~ 1 Gyr after the
first stars, on the metal-rich side of the knee, the decrease of
[/Fe] with increasing
metallicity is very well marked. In fact, the end points of the
evolution in each of the dSphs investigated has a surprisingly low
[
/Fe], see
Fig. 11.
A natural explanation of these low ratios could involve a sudden decrease
of star formation, that would make enrichment by massive stars
inefficient and leave SNIa to drive the chemical evolution. This
sudden drop in star formation could be the natural result of galactic
winds which can have a scientific impact on dwarf galaxies with
relatively shallow potential wells (see
Section 5) or perhaps tidal stripping.
In this case, one would expect the metal-rich and low
[
/Fe]
populations to be predominantly young, corresponding to the residual
star formation after the sudden decrease. However, the current
age-determinations for individual giants in these systems are not
accurate enough to probe this hypothesis (e.g.,
Battaglia et
al. 2006).
4.1.2 SODIUM AND NICKEL Another example of the low impact of massive-stars on the chemical enrichment of dSphs is given by sodium. Fig. 12 shows the compilation of dSphs stars compared to the evolution of Na in the MW. According to stellar current models, Na is mostly produced in massive stars (during hydrostatic burning) with a metallicity-dependent yield. The abundance of Na in metal-poor dSph stars is apparently not different to the MW halo stars at the same [Fe/H], but its abundance at later stages in the evolution is distinct from the MW, dSph producing (or keeping) too little Na to keep on the MW trend above [Fe/H] > -1.
![]() |
Figure 12. Sodium (above) and nickel (below) in the same four dSphs as in Fig. 11, compared to the MW. Sgr (red: Sbordone et al. 2007, Monaco et al. 2005, McWilliam & Smecker-Hane 2005), Fnx (blue: Letarte 2007, Shetrone et al. 2003), Scl (green: Hill et al. in prep; Shetrone et al. 2003, Geisler et al. 2005) and Carina (magenta: Koch et al. 2008a, Shetrone et al. 2003). Open symbols refer to single-slit spectroscopy measurements, while filled circles refer to multi-object spectroscopy. The small black symbols are a compilation of the MW disk and halo star abundances, from Venn et al. (2004a). |
Sodium and nickel under-abundances have also been remarked upon by
Nissen & Schuster
(1997,
2009)
in a fraction of halo stars which also
display low abundances,
thereby producing a [Na/Fe] - [Ni/Fe]
correlation. This correlation is tentatively explained as the common
sensitivity of both elements to neutron-excesses in supernovae.
Fnx is the most striking example that seems to
follow the same slope as the Na-Ni relationship in the MW, but extending
the trend to much lower [Na/Fe] and [Ni/Fe] values
(Letarte 2007),
see Fig. 12.
Nickel, unlike sodium, is also largely produced in SNIa
(Tsujimoto et
al. 1995),
so the Ni-Na relation can in theory be modified by SNIa
nucleosynthesis, especially in the metal-rich populations of dwarfs
where the low [
/Fe]
ratios point towards a strong SNIa contribution.
4.1.3 NEUTRON-CAPTURE ELEMENTS
Despite their complicated nucleosynthetic origin, heavy neutron
capture elements can provide useful insight into the chemical
evolution of galaxies. Nuclei heavier than Z ~ 30 are produced by
adding neutrons to iron (and other iron-peak) nuclei. Depending on the
rate (relative to
decay) at
which these captures occur, and
therefore on the neutron densities in the medium, the processes are
called either slow or rapid (s- or r-) process.
The s-process is well
constrained to occur in low to intermediate-mass (1-4
M
)
thermally pulsating AGB stars (see
Travaglio et
al. 2004
and references therein), and therefore provide a contribution to
chemical enrichment that is delayed by ~ 100-300 Myrs from the
time that the stars were born. Thus s-process elements can in
principle be used to probe star formation on similar timescales to
[
/Fe].
The r-process production site is clearly associated
with massive star nucleosynthesis. The most plausible candidate being
SNII, although the exact mechanism to provide the very large neutron
densities needed is still under debate, (e.g.,
Sneden, Cowan &
Gallino 2008
and references therein). This means that r-process elements should
contribute to the chemical enrichment of a galaxy with very little, if
any, delay. Obviously they need pre-existing Fe-peak seeds and are
therefore not primary elements such as
elements. One
complication arises from the fact that most neutron-capture elements
(through their multiple isotopes) can be produced by either the s- or
the r- process, such as yttrium (Y), barium (Ba) or lanthanum (La).
Among the few exceptions is europium (Eu), which is almost exclusively
an r-process product.
Fig. 13 compares Ba and Eu abundances in four
dSph galaxies and in the MW. At first glance, the Eu evolution
in dSph galaxies resembles that of their respective
-elements
(see Fig. 11), as expected for an
r-process originating
in massive stars. In the MW, the Ba and Y are dominated by the
r-process for [Fe/H]
-2.0 (e.g.,
Simmerer et
al. 2004,
Johnson &
Bolte 2002),
while the s-process dominates at higher metallicities (e.g., more
than 80% of the solar Ba is of s-process origin).
![]() |
Figure 13. Neutron-capture elements Y, Ba & Eu in the same four dSphs as in Fig. 11, compared to the MW. Sgr (red: Sbordone et al. 2007, Monaco et al. 2005, McWilliam & Smecker-Hane 2005), Fnx (blue: Letarte 2007, Shetrone et al. 2003), Scl (green: Hill et al. in prep; Shetrone et al. 2003, Geisler et al. 2005) and Carina (magenta: Koch et al. 2008a, Shetrone et al. 2003). Open symbols refer to single-slit spectroscopy measurements, while filled circles refer to multi-object spectroscopy. The small black symbols are a compilation of the MW disk and halo star abundances, from Venn et al. (2004a). |
At early times (at [Fe/H] < -1) there seems to be little difference
between the various dSphs, and the MW halo in
Fig 13. However, there is a hint that at the
lowest metallicities
([Fe/H] < -1.8), [Ba/Fe] increases in scatter and starts to turn
down. This hint is confirmed in the plot in
Section 4.2 which
includes other dSphs, although from much smaller samples
(Shetrone,
Côté & Sargent 2001,
Fulbright, Rich
& Castro 2004,
Aoki et al. 2009).
In fact, this scatter and downturn of [Ba/Fe] is a well known feature in
the MW halo
(François et
al. 2007,
Barklem et
al. 2005
and references therein), where it occurs at much lower
metallicities ([Fe/H] < -3.0). So far we have extremely low number
statistics for dSphs and these results need to be confirmed in larger
samples of low-metallicity stars. These low r-process values at
higher [Fe/H] than in the Galactic halo would either mean that the
dwarf galaxies enriched faster than the halo at the earliest times or
that the site for the r-process is less common (or less efficient) in
dSphs. The r-process elements are clearly useful tracers of early
time scales, because unlike the
-elements (in the halo
and in dSphs)
they show significant scatter in the lowest metallicity stars. The
r-process is thus produced in much rarer events than the
-elements and so it can
be a much finer tracer of time scales and enrichment (and mixing) processes.
The ratio of [Ba/Eu], shown in Fig. 14,
indicates the
fraction of Ba produced by the s-process to that produced by the
r-process. In dSphs, as in the MW, the early evolution of all
neutron-capture elements is dominated by the r-process (this was
already noted by
Shetrone,
Côté & Sargent 2001,
Shetrone et
al. 2003).
In each system, however,
the low and intermediate mass AGB stars contribute s-process
elements, that soon start to dominate the Ba (and other neutron capture
elements) production. The metallicity of this switch from r-
to s-process ([Fe/H] ~ -1.8, the same as the
[/Fe] knee)
is only somewhat constrained in the Scl dSph.
This turnover needs to be better constrained in
Scl and even more so in other galaxies to provide
timing constraints on the chemical enrichment rate. It could reveal the
metallicity reached by the system at the time when the s-process
produced in AGBs starts to contribute.
![]() |
Figure 14. Ratios of r- to s- process element production in the same four dSphs as in Fig. 11, compared to the MW. Sgr (red: Sbordone et al. 2007, Monaco et al. 2005, McWilliam & Smecker-Hane 2005), Fnx (blue: Letarte 2007, Shetrone et al. 2003), Scl (green: Hill et al. in prep; Shetrone et al. 2003, Geisler et al. 2005) and Carina (magenta: Koch et al. 2008a, Shetrone et al. 2003). Open symbols refer to single-slit spectroscopy measurements, while filled circles refer to multi-object spectroscopy. The small black symbols are a compilation of the MW disk and halo star abundances, from Venn et al. (2004a). |
For the more metal-rich stars ([Fe/H] > -1) there is also a distinctive behaviour of [Ba/Fe] in dSphs (Fig. 13). In the Scl dSph the [Ba/Fe] values never leave the MW trend, but this galaxy also has almost no stars more metal-rich than [Fe/H] < -1. Fnx, on the other hand, and to a lesser extent Sgr, display large excesses of barium for [Fe/H] > -1. This is now barium produced by the s-process, and it shows the clear dominance of the s-process at late times in dSphs.
Fig. 15 shows the trends of
[Fe/],
and [Ba/
] against
[
/H]. The fact that
[
/H] keeps increasing
significantly after the knee in the Scl dSph
demonstrates that even in this system which has no significant
intermediate-age population, there was still ongoing star formation
contributing
enrichment
from massive stars well after SNIa started contributing.
This is also confirmed by the presence of stars, which have
[Fe/H] > -1.8 (the knee), and were therefore formed after SNIa
started exploding. In Fnx or Sgr, the very flat, extended and high
[Fe/
] plateau also shows
that massive stars have kept feeding
the chemical enrichment all along the evolution, even though they do
not dominate the Fe enrichment. As for the s-process, the widely
different behaviour of Scl, Fnx and Sgr and the MW are even
more striking viewed in this representation than they were in
Fig. 13, illustrating the total disconnect of
massive stars
nucleosynthesis to Ba, and the strong influence of AGB stars at a time
when massive stars do not drive the metallicity evolution anymore.
![]() |
Figure 15. Trends of iron and neutron-capture elements as a function of (PUT IN THE ALPHA SYMBOL HERE) elements in the same four dSphs as in Fig. 11, compared to the MW. Sgr (red: Sbordone et al. 2007, Monaco et al. 2005, McWilliam & Smecker-Hane 2005), Fnx (blue: Letarte 2007, Shetrone et al. 2003), Scl (green: Hill et al. in prep; Shetrone et al. 2003, Geisler et al. 2005) and Carina (magenta: Koch et al. 2008a, Shetrone et al. 2003). Open symbols refer to single-slit spectroscopy measurements, while filled circles refer to multi-object spectroscopy. The small black symbols are a compilation of the MW disk and halo star abundances, from Venn et al. (2004a). |
4.2. Ultra-Faint dwarf galaxies
Individual stars in the uFds that have recently been discovered around the MW have so far been little observed at high spectral resolution. This is probably due to the difficulty in confirming membership for the brighter stars in these systems. However, several groups are currently following up confirmed members (typically selected from lower resolution Ca II triplet observations) to derive abundances. So far, Koch et al. (2008b) have observed two RGB stars in Herc (MV ~ -6.6) and Frebel et al. (2009) are following up RGB stars in the even fainter uFds UMa II and Coma (both with, MV ~ -4.). The latter study confirms that uFds do contain very metal-poor stars, [Fe/H] < -3, (as found by Kirby et al. 2008), unlike the more luminous "classical" dSphs (Helmi et al. 2006). It also appears that these uFds extend the metallicity-luminosity relation down to the lower luminosities (Simon & Geha 2007), see Section 3.3.
The two stars in Herc seem to have particularly peculiar abundance
patterns, with high Mg and O abundances (hydrostatic burning in
massive stars), normal Ca, Ti abundances (explosive nucleosynthesis in
massive stars), and exceedingly unenriched in Ba
(Koch et
al. 2008b).
On the other hand, elemental ratios in the extremely metal-poor stars in
the two fainter dwarfs UMa II and Coma
(Frebel et
al. 2009)
are remarkably similar to the MW halo extremely
metal-poor stars. Fig. 16 compares Mg and Ba
measurements in faint dSphs, with all more luminous dSph stars that
have [Fe/H] -2
(including a new sample of 6 very metal-poor stars in Sextans by
Aoki et al. 2009)
and the MW. In fact, only Sextans seems
to have scattered and low [Mg/Fe] ratios, while other dSphs and uFds all
show similar [Mg/Fe] enhancements. The similarity between stars
with metallicities below [Fe/H]
-2.5 in the MW
and faint dwarfs is seen also in other light elements, such as Na, Sc,
Cr, Mn, Ni or Zn. This may also be true of more luminous dSphs, see
Section 4.1.
![]() |
Figure 16. Mg
( |
The overall similarity between all the most metal-poor stars for element ratios up to the iron-peak can be taken as an indication that star formation and metal-enrichment, even at the earliest times, and even in the smallest systems, has proceeded in a similar manner. This may lead to the net yield of the very first stars. The very low dispersion found in abundance ratios of these elements in Galactic extremely metal-poor stars (EMPS), down to metallicities of [Fe/H] ~ -4 came as a surprise (Cayrel et al. 2004): since it was thought that one or a few SN II were sufficient to enrich the gas to those metallicities, the expectation was that among EMPS the variety of metal-production sites (SN II of different masses) would appear as dispersed abundance ratios. We are now adding to this puzzle the fact that these well defined abundance ratios are also achieved by considerably smaller halos.
The only discrepancy among the most metal-poor stars concerns the r-process element Ba, that stands out below the MW halo distribution both for faint and somewhat more luminous dSph galaxies. The most extreme low Ba abundances are found so far in Herc (Koch et al. 2008b) and Draco (Fulbright, Rich & Castro 2004), where only upper limits were detected.
The dIs are all (except the SMC) located at rather large distances from the MW and so far, the only probes that could be used to derive chemical abundances in these objects were HII regions and a few super-giant stars. Both types of probes allow a look-back time of at most a few 10 Myr, and this is the end-point of a Hubble-time's worth of chemical evolution for any galaxy. This limitation makes it difficult to gather relevant information to constrain the chemical enrichment over time in these systems. However, abundances in HII regions and super-giants (see references in Table 1) are useful to understand how dIs fit in the general picture of dwarf galaxies, and how they compare to larger late-type galaxies. First, they give the present day metallicity of these systems, and all are more metal-poor than the MW disk young population, in agreement with the metallicity-luminosity relation (see for example van Zee & Haynes 2006 for a relation based on dIs within 5 Mpc), and range between 12 + log(O/H) ~ 8.1 (e.g., NGC 6822, IC 1613) to 12 + log(O/H) ~ 7.30 (Leo A), or [O/H] ~ -0.6 to -1.4. Both HII regions and super-giants typically agree on the oxygen abundances of the systems, within the respective measurement uncertainties (Venn et al. 2003, Kaufer et al. 2004), with little metallicity dispersion within a galaxy, and no spatial gradient (e.g., Kobulnicky & Skillman 1997, van Zee, Skillman & Haynes 2006, van Zee & Haynes 2006). This holds even in the most metal-poor galaxies, and suggests a very efficient mix of metals across the galaxy despite the clumpiness of the ISM and ongoing star-formation. The shear within these systems is expected to be very low, and this has been taken as an indication that mixing occurs in the gaseous hot phase, before the gas cools down to form new stars (e.g., van Zee, Skillman & Haynes 2006).
A to M type super-giants have a further interest as they provide the
present-day
[/Fe]
ratios in dIs, which are not accessible from HII regions
where typically only light elements (e.g., He, N, O, Ne, S, Ar) can be
measured, and and no iron (nor any
other element that would trace SNIa).
The first dI where abundances of stars were measured was of course the SMC in our backyard. The largest samples to date with abundances in SMC are of super-giants which can be found in Hill, Barbuy & Spite (1997, K-type stars), Luck et al. (1998 F-type stars) or Venn (1999 A-type stars). Similar studies in more distant dIs needed efficient spectrographs on 8-10m telescopes, and at the expense of observing for many hours a few stars detailed abundances have been observed in A-type super-giants out to distances of 1.3 Mpc. This work has been pioneered by K. Venn and collaborators using A-type stars (Venn et al. 2001, 2003, Kaufer et al. 2004) in NGC 6822, Sextans A and WLM. There has also been a more recent study using M-type stars in IC 1613 by Tautvaisiene et al. (2007).
Fig. 17 illustrates the observed low
[/Fe] in
these systems, and compares them to the observed trends of older
populations (RGB stars) in dSph galaxies, as defined in
Figs. 16 & 11. These
low [
/Fe] are
expected in galaxies that have formed stars over a long period of
time, however; they clearly occur at much lower metallicities than in
larger systems such as the MW or the LMC, pointing towards an
inefficient metal-enrichment of the galaxy (low star formation and/or
metal-losses through winds). It is interesting to see in
Fig. 17 how dIs actually prolong the trends of
dSph galaxies, not only for
elements but also for
neutron-capture
elements. From these diagnostics, dSphs are entirely consistent with
being dIs that lost their gas at a late stage of their evolution. The
Fnx dSph and the SMC, which are both dominated by intermediate-age
populations, are also quite similar in their chemical enrichment,
except that Fnx ran out of gas (or lost its gas) and stopped star
formation about 108yr ago.
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Figure 17. Mg
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