Scott Trager 4 June 2007

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Dwarf galaxies in the Local Group

Scott Trager

4 June 2007

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Why should we care?

• In a hierarchical galaxy formation scenario, the smallest galaxies should form first and then merge to form larger galaxies

• May have contained the first stars

• May be shredded to make galaxy halos

• Does this happen? Do we see that dwarfs

are the first galaxies? Is our halo made up

of shredded dwarfs?

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Types of dwarf galaxies

• Two basic types:

• dwarf irregular (dI or dIrr) galaxies

• ongoing star formation, gas, dust

• dwarf spheroidal (dSph) galaxies

• no current star formation, no gas, little

(no?) dust

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• dIrr: GR8 in Virgo

• outlying member of LG

• note bright HII regions

Source: Diedre Hunter, Lowell Obs.

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• The Fornax dSph

• no ongoing star formation: no HII regions

Source: Skyview

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• Additional types:

• Blue Compact Dwarfs (BCDs)

• massive extension of dIrr sequence

• Dwarf ellipticals (dEs)

• massive extension of dSph sequence

• not low-mass ellipticals!

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Kinematic properties of dwarfs

Dwarfs follow a “fundamental line” (Dekel & Woo 2003): surface brightness, circular velocity, and

metallicity all correlated with (stellar) mass

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Integrated properties

• dIrrs tend to be blue and

(typically) brighter than dSphs

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July 25, 1998 17:37 Annual Reviews AR062-11

LOCAL GROUP DWARFS

459

Figure 5 The M_V₀-(B−V)₀ color-magnitude diagram of Local Group galaxies based on the data from Tables 2 and 3. The diagonal dashed lineseparates galaxies classified as Spirals, or Irregular systems (filled squares; Table 1), and dSph or Elliptical systems (open circles). Five “transition”

objects, Antlia, LGS 3, Phoenix, Pegasus, and DDO 210, are plotted as filled triangles; NGC 205 is plotted as a filled circle.

Demers et al (1994b, 1995) carefully searched for substructure in a number of dSph systems but found weak evidence for such structure only in Ursa Minor (Olszewski & Aaronson 1985).

Only three Local Group dwarfs contain nuclei: NGC 205, Sagittarius, and M32. The latter is now widely believed to contain a massive central black hole (Kormendy & Richstone 1995). The nucleus of NGC 205 is extremely blue (Price & Grasdalen 1983, Lee 1996), dynamically colder than the surrounding galaxy envelope (Carter & Sadler 1990), and has a spectrum dominated by young stars (Bica et al 1990, Jones et al 1996). The existence of a nucleus of Sagittarius—the globular cluster M54—is somewhat controversial. Da Costa

& Armandroff (1995) argued that the velocity dispersion and metallicity of M54 are incompatible with its identification as a normal dSph nucleus. However, M54 is the second most luminous globular cluster in the entire Milky Way, nearly as luminous as the nucleus of NGC 205 (Peterson 1993); exhibits an internal abundance dispersion (Sarajedini & Layden 1995); and is located close to the center of symmetry of Sagittarius (Ibata et al 1994, 1997). Such an unusual object seems unlikely to have been merely an isolated globular cluster in a dSph galaxy such as Sagittarius.

Annu. Rev. Astro. Astrophys. 1998.36:435-506. Downloaded from arjournals.annualreviews.org by University of Groningen on 03/09/05. For personal use only.

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• (LG) dwarfs are dark matter

dominated

(Aaronson 1983), particularly at the low-mass end

4 Strigari et al.

Fig. 1.— The likelihood functions for the mass within 0.6 kpc for the nine dSphs, normalized to unity at the peak.

Sagittarius is further complicated by the fact that it is experiencing tidal interactions with the MW (Ibata et al.

1997; Majewski et al. 2003), so a mass estimate from the Jeans equation is not necessarily reliable. We caution that in this case the mass we determine is likely only an approximation to the total mass of the system.

We determine the likelihoods by marginalizing over the following ranges of the velocity anisotropy, inner and outer slopes: −10 < β₀ < 1, −10 < β∞ < 1, 0.1 < r_β < 10 kpc, 0.7 < γ < 1.2, and 2 < δ < 3.

As discussed above, these ranges for the asymptotic inner and outer slopes are appropriate because we are con- sidering CDM halos. It is important to emphasize that these ranges are theoretically motivated and that observations alone do not demand such restrictive choices.

It is possible to fit all of the dSphs at present with a constant density cores with scale-lengths ∼ 100 pc (Strigari et al. 2006; Gilmore et al. 2007), although the data by no means demand such a situation. Though we consider inner and outer slopes in the ranges quoted above, our results are not strongly affected if we widen these intervals. For example, we find that if we allow the inner slope to vary down to γ = 0, the widths of the likelihoods are only changed by ∼10%. This reflects the fact that there is a negligible degeneracy between M_0.6 and the inner and outer slopes.

We are left to determine the regions of ρs and rs parameter space to marginalize over. In all dSphs, there is a degeneracy in this parameter space, telling us that it is not possible to derive an upper limit on this pair of parameters from the data alone (Strigari et al. 2006).

While this degeneracy is not important when determining constraints on M0.6, it is the primary obstacle in determining V_max. From the fits we present below, we find that the lowest rs value that provides an acceptable fit is ∼ 0.1 kpc, and we use this as the lower limit in all cases. In our fiducial mass models, we conservatively re- strict the maximum value ofrs using the known distance to each dSph. In this case, we use 0.1 < rs < D/2, where D is the distance to the dSph.

In Figure 1 we show the M0.6 likelihood functions for

Fig. 2.— The mass within 0.6 kpc (upper) and the mass-to-light ratios within the King tidal radius (lower) for the Milky Way dSphs as a function of dwarf luminosity. The error-bars here are defined as the locations where the likelihoods fall to 40% of the peak values (corresponding to ∼1σ errors).

all of the dSphs. As is shown, we obtain strong constraints onM_0.6 in all cases except Sagittarius, for which we use only a central velocity dispersion. Table 1 sum- marizes the best fitting M0.6 values for each dwarf. The quoted errors correspond to the points where the likelihood falls to 10% of its peak value. The upper panel of Figure 2 shows M_0.6 values for each dwarf as a function of luminosity.

The maximum likelihood method also allows us to constrain the mass at other radii spanned by the stellar distribution. The sixth column of Table 1 provides the integrated mass within each dwarf’s King tidal radius. This radius roughly corresponds to the largest radius where a reasonable mass constraint is possible. As expected, the mass within rt is not as well determined as the mass within 2rking. From these masses we are able to determine the mass-to-light ratios withinr_t, which we present in the seventh column of Table 1. In the bottom panel of Figure 2, we show mass-to-light ratios within r_t as a function of dwarf luminosity. We see the standard result that the observable mass-to-light ratio increases with de- creasing luminosity (Mateo 1998). Note, however, that our results are inconsistent with the idea that all of the dwarfs have the same integrated mass within their stellar extent. We note that for Sagittarius, we can only obtain a lower limit on the total mass-to-light ratio.

The last two columns in Table 1 list constraints on Vmax for the dSphs. Column 8 shows results for an analysis with limits onr_s as described above. In this case, the integrated mass within the stellar radius is constrained by the velocity dispersion data, but the halo rotation velocity curve, V_c(r), can continue to rise as r increases beyond the stellar radius in an unconstrained manner.

The result is that the velocity dispersion data alone pro- vide only a lower limit on V_max.

Stronger constraints on Vmax can be obtained if we limit the range of rs by imposing a cosmology-dependent prior on the dark matter mass profile. CDM simula- tions have shown that there is a correlation between

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July 25, 1998 17:37 Annual Reviews AR062-11

490 ^MATEO

Figure 9 Kinematically determined mass-to-light ratios of local group dwarfs as a function of luminosity. Top panel: log(M/L)0 from Table 4 vs MV. Filled squaresare for dSph or dSph/Irr systems for which masses were determined from the central velocity dispersions, while theopen squaresrepresent Irr systems that have masses derived here from HI rotation curves. See Table 4 for details or the original sources to obtain definitive kinematic mass estimates for these galaxies.

Sagittarius is denoted as an open circle. Bottom panel: log(M/L)_tot from Table 4 vs M_V; the symbolsare the same as in thetop panel. In each panel I have also plotted the function logM/L = 2.5+10⁷/(L/L_!)as adashed line.

Annu. Rev. Astro. Astrophys. 1998.36:435-506. Downloaded from arjournals.annualreviews.org by University of Groningen on 03/09/05. For personal use only. Mateo (1998)Strigari et al. (2007)

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Spatial distribution of LG dwarfs

• dSph’s tend to live close to a giant galaxy (MW, M31)

• dIrr’s tend to be

distributed throughout the LG

2 Grebel

Group reviews have appeared in the past few years, including Mateo (1998), Grebel (1997, 1999, 2000a), and the very detailed recent reviews by van den Bergh (1999, 2000). Reviews dealing with dark matter in Local Group dwarf spheroidals include Mateo (1997) and Olszewski (1998).

Sextans N3109/

Group

NGC 6822

And II IC 10 M33

M31 ^{NGC 185}

WLM

Car

For

Dra UMi Sex

Cet

And III And I

Leo II

Leo A Tuc

Milky Way 500 kpc

500 kpc

NGC 3109 Antlia

Sex B

Sex A

And V

LGS 3

NGC 147 NGC 205 Cas dSph

And VI Peg dI

IC 1613

Sgr Scl SMC & LMC

Phe

Leo I DDO 210

SagDIG

Sgr Group

Local Group

1 Mpc

2Mpc

Figure 1. A scaled 3-D representation of the Local Group (LG). The dashed ellipsoid marks a radius of 1 Mpc around the LG barycenter (assumed to be at 462 kpc toward l = 121.7 and b = −21.3 following Courteau & van den Bergh 1999). Distances of galaxies from the the arbitrarily chosen plane through the Milky Way are indicated by solid lines (above the plane) and dotted lines (below). Morphological segregation is evident: The dEs and gas-deficient dSphs (light symbols) are closely concentrated around the large spirals (open symbols).

DSph/dIrr transition types (e.g., Pegasus, LGS 3, Phoenix) tend to be somewhat more distant. Most dIrrs (dark symbols) are fairly isolated and located at larger distances. Also indicated are the locations of two nearby groups.

2. Local Group galaxy content and distribution

The most massive and most luminous Local Group galaxies are the two spirals Milky Way and M31 (≈ 10¹²M!, MV ∼< −21 mag). The third, less luminous and less massive Local Group spiral M33 does not have any known compan- ions and belongs to the M31 subsystem. About two thirds of the Local Group galaxies are found within 300 kpc of the two spirals. The majority of these close

FIGURE 1. Morphological segregation in the number distribution N of different types of galaxies in the Local Group (solid histograms) and in the M81 and Centaurus A groups (hashed histograms) as a function of distance D to the closest massive primary (updated version of Fig. 1 in Grebel 2004).

1.4. Direct evidence for harassment and accretion

If accretion is the primary mechanism for the growth and evolution of massive galaxies as predicted by cosmological models, then we should be able to find evidence for these processes in our immediate neighborhood. (1) The study of the structural properties of nearby galaxies can reveal whether external tidal forces are distorting them. (2) The detection of extratidal stars and streams around and within massive galaxies is evidence for ongoing harassment and accretion events. (3) The stellar content, population properties, and chemistry of nearby galaxies allow us to constrain to what extent these kinds of objects could have contributed as building blocks to more massive galaxies.

The clearest evidence for ongoing accretion are the extended tidal stream of the Sagit- tarius dSph galaxy (Ibata, Gilmore, & Irwin 1994), which has now been traced around the entire Milky Way using M giants identified in the Two Micron All-Sky Survey (2MASS) (Majewski et al. 2003), and the giant stream of metal-rich giants in the halo of M31 (Ibata et al. 2001; Ferguson et al. 2002). Additional stellar overdensities have been detected in the Milky Way using various photometric data sets including 2MASS and the Sloan Digital Sky Survey (SDSS): The Monoceros feature (Newberg et al. 2002;

Yanny et al. 2003), which may be a tidal tail connected with the Canis Major overdensity (Martin et al. 2004a). The interpretation of Canis Major is disputed; suggestions include that it is part of the Galactic warp or flare (Momany et al. 2004) or indeed the center of another possibly disrupted dSph within the Milky Way (Martin et al. 2004a, 2004b).

Additional Galactic stellar overdensities have been identified (for instance, Triangulum-

Grebel (2000)Grebel (2005)

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-1 -1

-2 -2 -2 -2

0 0

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& field clusters

spread spread const enrich.

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[Gyr]

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& field WR

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15 15

10 5 0 10 5 0 10 5 0

15

1710 kpc 1800 kpc

1610 kpc 1700 kpc 1590 kpc

590 kpc 440 kpc 870 kpc

NGC 6822

650 kpc

0 0

0 5 15 5 15 5 15 5 15 5

15 10 10 10 10 10

15 10 5 0 15 10

Star Formation Histories of Irregular Galaxies

5 0 15 10 5 0 15 5 0 15 10 5 0

430 kpc 470 kpc

IC 10

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WLM

770 kpc

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8 GCs

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(Distances from barycenter of Local Group) OB ass. = OB associations, YSG= yellow supergiants, TRGB= tip of the red giant branch, HII = HII regions, WR= W-R stars, Cep= Cepheids, BL= blue loop stars

Grebel 1998

Figure5:PopulationboxesforLGdIrr(Romanfonts)andanddIrr/dSphgalaxies(italics).

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15 NGC 205 And II

And III And I

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~100 kpc

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123 kpc 215 kpc 879 kpc

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15 10 5 0 15 10 5 0 15 10 5 0 15 10 5 0 15 10 5 0

Fornax

Sculptor

Star Formation Histories of dSph / dE Galaxies

Sextans Draco

Ursa Minor

87 kpc

71 kpc 80 kpc

25 kpc

Carina

MW ^MW ^MW 86 kpc ^MW

MW MW MW MW

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?

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Grebel 1998

Figure6:PopulationboxesforLGdSph(Romanfonts)anddEgalaxies(italics).

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• No two LG dwarfs have the same star formation history (Mateo 1998)

• but fall into broad classes

• All LG dwarfs have old populations (Grebel &

Gallagher 2004)

Lessons learned from stellar populations

Reionization (z ~ 20, WMAP)

Reionization (z ~ 6.4, SDSS) Big Bang(WMAP)

5 6 7 8 9 10 11 12 14

4 13

Time before present (Gyr)

z = 0.5 z = 1 z = 2 z = 3

dSphs with episodic star formation (Carina)

"Young" dSphs; transition types

"Intermediate−age" dSphs

"Old" dSphs (Draco) (Fornax, Phoenix)

(Leo I)

FIGURE 2. Bar diagram indicating the approximate duration of star formation episodes in low-mass galaxies (∼10⁷ M"). The approximate beginning and end of the re-ionization epoch are indicated (based on results from WMAP and from the Sloan Digital Sky Survey). The predicted cessation of star formation in low-mass galaxies, esp. in dSphs, is not observed. For more details, see also Grebel & Gallagher (2004).

These are testable predictions that can be investigated by exploiting the fossil stellar record in nearby galaxies. The LG is an ideal target since here the oldest populations are resolved and can be accessed with HST and increasingly also with large ground-based telescopes operating with high angular resolution. This requires age dating of old populations. The most accurate ages can be obtained for resolved stellar populations. For old populations the most age-sensitive feature is the old main-sequence turn-off (MSTO), where via differential age-dating techniques internal accuracies of less than 1 Gyr can be obtained. Comparison objects are usually ancient Galactic globular clusters of the same metallicity as the target population. Absolute ages are more difficult to determine since here one needs to rely on isochrones, which makes the resulting ages model-dependent.

But also differential techniques have a number of drawbacks: They require very deep, high-quality photometry reaching at least 2 mag below the MSTO, which is a challenge in more distant galaxies. They require stellar populations sufficiently numerous to produce a measurable MSTO, which necessarily limits us to Population II stars (note that to date not a single Population III star candidate has ever been detected beyond the Milky Way). Ideally, one wishes to compare populations with the same [!/Fe] ratio. (Estab- lishing to what extent this condition is met and what the consequences are will be an important area for large ground-based telescopes in the coming years.) Other effects such as diffusion may further affect relative measurements; see Chaboyer’s contribution in these proceedings for details.

For a more detailed discussion of the results of differential age dating as applied

Grebel (2005)

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Chemistry of LG dSphs

• Big problem: dSph

abundance patterns do not match MW halo star pattern (Shetrone et al.

2003; Tolstoy et al. 2003)

• can’t form MW halo from present-day

dwarfs

36 The Messenger 123 – March 2006

vations of stars in globular clusters with high-resolution abundances. Assum- ing sufficient signal-to-noise spectra (S/N > 10) it provides [Fe/H] within typical (internal) errors of ± 0.1 dex, and also the radial velocity of each star with ± 2 km/s accuracy. These accuracies are well suit- ed for ‘quick look’ surveys of the resolved stellar population of a galaxy. In the DART project these CaII triplet measurements are complementary to the high-resolution observations made in the centre of each dSph. In the low-resolution observations a much larger area is surveyed and we can assess how representative the detailed study is of the stellar population of the whole galaxy.

Our first VLT/FLAMES results (Tolstoy et al. 2004), were based upon CaII triplet measurements, which clearly showed that Sculptor dSph contains two distinct stellar components with different spatial, kinematic and abundance properties (see Figure 5). The upper panel shows the VLT/FLAMES spectroscopic measurements of [Fe/H] for 307 probable velocity members of Sculptor (with S/N > 10).

We see a clear trend of metallicity with radius. The lower panel shows v_hel as a function of elliptical radius for all stars

satisfying S/N > 10. Likely Sculptor members are clearly seen clustered around the systemic velocity of 110 km/s. The stars which are potential members are plotted as red stars ([Fe/H] > −1.7) and blue circles ([Fe/H] < −1.7), while the green crosses are assumed to be non- members. There appears to be a metal- rich, −0.9 > [Fe/H] > −1.7, and a metal- poor, −1.7 > [Fe/H] > −2.8, component.

The metal-rich component is more centrally concentrated than the metal-poor, and on average appears to have a lower velocity dispersion, S_metal−rich = 7 ± 1 km/s, whereas S_metal−poor = 11 ± 1 km/s. A similar effect is seen in Fornax, where the metal-rich stars are centrally concentrated, and the metal-poor stars appear more uniformly and diffusely distributed.

It is clear from the histogram of [Fe/H]

measurements that both Sculptor and Fornax lack a low metallicity tail (see Fig- ure 6). In Figure 6 are plotted in the left- hand histogram distributions of [Fe/H]

for Sculptor dSph: the 91 stars within the central, r < 0.2 degree region (solid black line); and the 216 stars beyond r > 0.2 de- grees (dashed red line). In Figure 6 in the right-hand histogram is the [Fe/H] distribution for the Fornax dSph: the 332 stars

within the central, r < 0.5 degree region (solid black line); and the 229 stars beyond r > 0.5 (dashed red line). Clearly the distributions are very different. Most noticeably Fornax has a substantial

‘metal-rich’ population. The lowest metallicity star in our combined sample of more than 850 stars for both galaxies is [Fe/H] = −2.7. We find a similar lack of low-metallicity stars in Sextans and Ca- rina. Although it is difficult to make an accurate comparison with the Galactic surveys, where the completeness can be hard to quantify, there appears to be a significantly different distribution between all the dSph and the (metal-poor) halo of the Milky Way. It can be seen that there is a clear difference in the distribution of the metal-rich and metal-poor stars in both Fornax and Sculptor.

Future work

There are indications that the presence of two distinct populations is a common feature of dSph galaxies. Our preliminary analysis of Horizontal Branch stars, v_hel and [Fe/H] measurements in the other galaxies in our sample (Fornax and Sextans dSph; Battaglia et al. 2006, in prep.) also show similar characteristics to Sculptor, especially in the most metal- poor component. Pure radial-velocity studies (e.g., Wilkinson et al. 2004) have also considered the possibility that kinematically distinct components exist in Ursa Minor, Draco and Sextans dSph galaxies. Interestingly, the Carina dSph appears to go counter to this trend, and another VLT/FLAMES study finds no ob- vious evidence for more than one component, or even a gradient within Carina dSph (Koch et al. 2006).

What mechanism could create two or more distinct ancient stellar components in a small dwarf spheroidal galaxy? A simple possibility is that the formation of these dSph galaxies began with an initial burst of star formation, resulting in a stellar population with a mean [Fe/H] ≤ −2.

Subsequent supernova explosions from this initial episode could have been sufficient to cause gas (and metal) loss such that star formation was inhibited until the remaining gas could sink deeper into the centre. Thus the subsequent gen- eration(s) of stars would form in a region Reports from Observers Tolstoy E. et al., The DART Large Programme – A Close Look at Nearby Galaxies

Halo Disc

Scl

[/Fe]

–4 –3 –2 –1 0

–0.4 –0.2 0 0.2 0.4 0.6

[FeI/H]

Halo Disc

Fnx

[/Fe]

–4 –3 –2 –1 0

–0.4 –0.2 0 0.2 0.4 0.6

[FeI/H]

Figure 4: The A-element abundances from high- resolution spectroscopy for stars in the Sculp- tor dSph (upper panel, Hill et al. 2006 in prep) and also preliminary results for the Fornax dSph (lower panel, Letarte et al. 2006a, in prep.). The open squares are UVES measurements of individual stars

(Shetrone et al. 2003; Geisler et al. 2005; Venn et al. 2006, in prep.). Galactic stars coming from various literature sources (see Venn et al. 2004 for refer- ences) are shown for comparison in both panels as black dots and labelled as disc and halo components.

Tolstoy et al. (2006)

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• Not so bad for r- and s- process elements?

• Well, [Ba/Y] (low- vs.

high-mass s-process) not the same

• low-metallicity AGB stars (e.g., Letarte 2007)?

this in Shetrone et al. (2003) after observing that 50% of our dSph stars with ½Fe=H" # $1:0 show the Na-Ni relationship.

The low [!/Fe] ratios in the NS97 sample, however, do indicate something odd in the formation of those stars (discussed further inx6.1).

A slight and positive correlation between Na and Ni is a natural result of nucleosynthesis in massive stars (see Clayton 1983,x5.6). As Timmes et al. (1995) suggest, Na is made by massive stars and delivered to the ISM in the SN II events.

The amount of Na produced is controlled by the neutron ex-

cess (primarily through the²²Ne content)¹⁰during hydrostatic C burning. While ²³Na is naturally produced during hydrostatic C burning in the core, the amount of ²³Na produced relative to ²⁴Mg ranges from 1:2 to 1:5, depending on the temperature for C burning (related to the mass of the star). But what is noteworthy is that²³Na is the only stable neutron-rich

Fig.4.—Ratios of [ Y/ Eu], [ Ba /Eu], [ La /Eu], and [ Ba /Y]. These ratios are used to examines-process enrichments. The [La /Eu] ratios suggest there is no significants-process enrichment until½Fe=H" # $1:8 in both the Galactic and dSph stars ([ Ba/ Eu] should also show this but may be compromised as discussed in the text). That [ Y/ Eu] is clearly lower in the dSph stars at all metallicities suggests differences in both ther- ands-process contributions (see text). The pure r-process estimates from solar system abundances are shown from Arlandi et al. (1999;dashed line) and Burris et al. (2000;dotted line). Same symbols as in Figs. 1 and 2.

10The²³Na yield may depend on the initial heavy element abundance in a star (e.g., Arnett 1971) because the²²Neproductionin a star is not strongly metallicity dependent (Arnould & Norgaard 1978).

VENN ET AL.

1186 Vol. 128

Venn et al. (2004)

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Models of dwarf galaxy formation

• Dekel & Silk (1986)

• galaxies with V c <120 km/s cannot retain gas due to supernova feedback

• dichotomy of properties at this velocity:

• low stellar mass, diffuse (low SB) dwarfs

• high stellar mass, compact (high SB)

“normal” galaxies

1986ApJ...303...39D

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• Ferrara & Tolstoy (2000) have extended this idea to show that galaxies with M<5x10 ⁶ M ^⊙☉ are

“blown away”

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DM fraction

Baryon mass

(18)

• Are there galaxies with masses below the “blow away” mass?

• Martin et al. (2007) have measured

velocity dispersion in some of the

“new” (SDSS) dwarfs

• maybe Wilman1? if so, how did it

survive?

14 N. F. Martin et al.

Figure 11. Comparison of the mass-to-light ratios (M/L) and mass estimates (M) for the faintest known dwarf galaxies with central velocity dispersion estimates (filled stars). In all cases, the mass estimates were derived using equation (3).

For Draco and Ursa Minor, we use the core radii and velocity dispersions quoted in Irwin & Hatzidimitriou (1995) and Armandroff, Olszewski & Pryor (1995) while we use those quoted in Chapman et al. (2005) for AndIX. For UMaI, the hollow star represent the 99% confidence higher mass limit obtained from the cold component they it may harbor (see § 5) while for CVnI, the two hollow stars represent the two diverging estimates obtained using either the cold more metal-rich half of the sample or the hot metal-poor half Ibata et al. (2006).

(1998) relation between M/L and the luminosity of a dwarf galaxies in the Local Group. This is not surprising since this relation (M/L = 2.5 + 10⁷/(L/L!)) naturally assumes that dwarf galaxies have masses higher than 10⁷ M!. Even though the uncertainties are sizeable, Wil1 would need to be at least ten times more massive to follow the relation.

Keeping the same structural parameters, this would mean that the central velocity dispersion would have to be at least

√10 ∼ 3 times higher than the one we have measured, that is

∼ 13 km s⁻¹, which is hardly compatible with the DEIMOS observations. Does it mean that a numerous population of small satellites (such as the structurally similar Segue 1;

Belokurov et al. 2006c) residing in less massive dark matter halos than brighter dwarf galaxies (∼< 10⁷M^!) have until now eluded us? Or does it mean that Wil1 was once a more luminous and massive dwarf galaxy and that its outskirts

have been stripped out by tidal interaction with the Milky Way, leaving only its central population visible at present?

Firstly, it has to be noted that equation (2) and (3) are only valid for a system with a constant mass-to-light ratio.

The mass estimates determined here only correspond to the central regions of the satellites where stars can be used as tracers. Therefore, it does not rule out per se that Wil1 could be embedded in a dark matter halo that is as massive as brighter galaxies. However, the low central mass estimate of this dark matter halo compared to the brighter galaxies would still hint at a halo with a lower central density and still make Wil1 a peculiar object.

The heating of an initially colder central population could also produce the observed velocity dispersion. While the metallicity spread that is measured within Wil1 ar- gues against the simple disruption of a globular cluster, one could argue that Wil1 is the remnant of a CVnI-like structure with only the colder core population remaining after stripping of the hotter component by tidal interaction with the Milky Way. However, in such a scenario, a significant amount of material would have to have been stripped to explain the absence of an underlying hot component in the spectroscopic data, at odds with the photometric data that do not show tidal tails containing a significant part of the whole satellite (Willman et al. 2006). Moreover, according to Piatek & Pryor (1995) the influence of tidal interaction of the measured velocity dispersion becomes significant at high distance from the center of the satellite whereas for the Wil1 sample that extends to ∼ 2rhb, there is no visible increase of the dispersion with distance in Figure 10.

Although the number of stars is too low for a detailed analysis, this does not favor the presence of strong tidal tails.

Thus it would seem more likely that Wil1 is a highly dark matter dominated object although it resides in a much less massive dark matter halo than those of brighter dwarf galaxies such as Boo. The small systemic velocity of this peculiar object (vgsr = 33.0 km s⁻¹) could mean it does not have a strongly radial orbit around the Milky Way, which could in turn explain why this object has survived until now.

Though dark matter halos are expected to form down to planet-mass structures (Diemand, Moore & Stadel 2005), a minimum mass is required to have a deep-enough potential to retain gas and eventually form stars. Therefore Wil1 could be of significant help in understanding how the lowest-mass systems form if it is confirmed to inhabit a low mass dark matter halo. Such a confirmation could come from the search for extra-tidal stars whose presence or absence would make it clearer if it is an unbound alignment of stars, a surviving core or a complete system. Besides, the presence of kinematically different populations in CVnI (Ibata et al. 2006) and perhaps UMaI yields significantly different mass estimates.

Plotting these mass estimates on Figure 11 (hollow stars) yields a significant spread in the mass-luminosity relation that should be taken as a warning against using M/L ratios when precise structural parameters are unknown. Moreover, it cannot be excluded at the moment that these satellites could have been strongly disrupted and influenced by tidal interaction or that they may have recently accreted some stellar material that did not have the time to relax in the gravitational potential of the dwarf. If this is the case, the simple application of equations (2) and (3) are not warranted and would lead to erroneous mass estimates.

c

! 0000 RAS, MNRAS 000, 000–000