Abundances from colourmagnitude diagrams

Currently, CMDs have been determined for at least seven (non-dwarf) elliptical and S0 galaxies: M32 (Grillmair et al. 1996), NGC 5128 (Rejkuba et al. 2005 and references therein), NGC 3379 (Gregg et al. 2004), NGC 3115, NGC 5102, NGC 404 (Schulte-Ladbeck etal. 2003), and Maffei 1 (Davidge & van den Bergh 2001). For all of these galaxies, CMDs have required HST imaging in the optical or near-infrared (NIR), except for Maffei 1, which was observed in the NIR with the CFHT adaptive-optics system. However, in only three cases - M32, NGC 3379 and NGC 5128 - have the CMDs reached significantly below the tip of the red giant branch (TRGB), as is required for reasonable abundance measurements (Figure 16.1, left). In the other cases, only the asymptotic giant branch (RGB) was detected (Maffei 1) or only stars less than one magnitude below the tip of the red giant branch (RGB) were detected. In NGC 3779, although the NICMOS F110W (J) and F160W (H) measurements extend to three magnitudes below the tip of the asymptotic giant branch (AGB), the tight spacing of isochrones of different metallicities and the s" -2 -

M32 Color Magnituded

Figure 16.2. Colour-magnitude diagrams of M32 (left) from Grillmair et al. (1996) and NGC 5128 (right) from Rejkuba et al. (2005). In the left panel, isochrones from Worthey (1994) have been over-plotted with metallicities in the range [Fe/H] = -1.2 to 0 dex in 0.2-dex steps for ages of 10 Gyr (dotted lines) and 15 Gyr (solid lines). Fiducial globular-cluster colour-magnitude sequences for NGC 6341, NGC 6752,47 Tuc, NGC 5927, and NGC 6528 (from left to right) have been over-plotted in the right panel.

F606W - F814W

F606W - F814W

Figure 16.2. Colour-magnitude diagrams of M32 (left) from Grillmair et al. (1996) and NGC 5128 (right) from Rejkuba et al. (2005). In the left panel, isochrones from Worthey (1994) have been over-plotted with metallicities in the range [Fe/H] = -1.2 to 0 dex in 0.2-dex steps for ages of 10 Gyr (dotted lines) and 15 Gyr (solid lines). Fiducial globular-cluster colour-magnitude sequences for NGC 6341, NGC 6752,47 Tuc, NGC 5927, and NGC 6528 (from left to right) have been over-plotted in the right panel.

co-mingling of the AGB and RGB over this magnitude range make precision estimates of stellar abundances very difficult, and only ranges of acceptable metallicities can be inferred, rather than accurate metallicity distributions (Figure 16.1, right) (Gregg et al. 2003).

In M32, Grillmair et al. (1996) obtained a CMD at a distance of r ~ 3re with WFPC2 on the HST. Because M31 is superimposed on M32, the resulting CMD was statistically cleaned of M31 stars by using a nearby M31 disc field. The resulting CMD reaches a depth just below the 'red clump', the pile-up of red horizontalbranch (helium-burning) stars about four magnitudes below the TRGB (Figure 16.2, left). There was no evidence of a blue horizontal branch (BHB) in this CMD, although the limiting magnitude in V prevented ruling out its presence completely. This absence has been confirmed in deeper WFPC2 (Worthey et al., in preparation) and ACS/HRC (Lauer et al., in preparation) observations. This is a crucial result, since BHB stars trace the most metal-poor stars in our Galaxy: at metallicities [Fe/H] < 1.5, helium-burning stars are dominantly BHB stars in Galactic globular clusters. Their lack in M32 suggests that this galaxy is deficient in metal-poor stars, which is confirmed by its metallicity distribution inferred from the distribution of RGB stars (Figure 16.3, left). Therefore M32 has a 'G-dwarf problem' (Worthey et al. 1996), for which the usual suggested solution is pre-enrichment of the galaxy (Tinsley 1980).

Figure 16.3. Metallicity distribution functions for M32 (left) from Grillmair et al. (1996) and NGC 5128 (right) from Rejkuba et al. (2005) inferred from comparing the distribution of RGB stars with isochrones. For M32, the isochrones used were 10 or 15 Gyr old (see legend); for NGC 5128, the isochrones used were 12 Gyr old. ForM32, closed-box chemical-evolution ('simple') models have been over-plotted with two different choices of the yield.

600 400 200 0

ars300

o 200 er

0 200 100

_. 1 1 1 1 1 1 1 1 . 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 l_ : 38 kpc field n n

: 21-31 kpc field BLgn

1 1 1 1 1 1 1 1.

¡W/tffln 1 , ,

Figure 16.3. Metallicity distribution functions for M32 (left) from Grillmair et al. (1996) and NGC 5128 (right) from Rejkuba et al. (2005) inferred from comparing the distribution of RGB stars with isochrones. For M32, the isochrones used were 10 or 15 Gyr old (see legend); for NGC 5128, the isochrones used were 12 Gyr old. ForM32, closed-box chemical-evolution ('simple') models have been over-plotted with two different choices of the yield.

For NGC 5128, Rejkuba et al. (2005), following earlier studies of inner fields by Harris & Harris (2000, 2002), obtained a CMD at a distance of r - 7re with the ACS/WFC on the HST. As for M32, Rejkuba et al. obtained a CMD down to the red clump (Figure 16.2, right) and found no evidence of a BHB. Again, it is difficult to rule out its presence completely, but the metallicity distribution inferred from the distribution of RGB stars (Figure 16.3, right) suggests that there are few (but some) very-metal-poor stars in NGC 5128.

By counting stars on the RGB between isochrones of different metallicities (but typically constant age), a metallicity distribution function (MDF) can be determined. The MDFs from M32 and NGC 5128 are shown in Figure 16.3. Both galaxies have a deficiency of metal-poor stars relative to a closed-box ('simple') chemical-evolution model, as discussed above. Galaxy M32 has a sharp peak in its MDF at [Fe/H] = -0.25, using isochrones from Worthey (1994). In Section 3.2 below I will compare this inferred MDF with results from stellar-population analyses. The MDF of NGC 5128 has a significant radial gradient and peaks at [Fe/H] ~ -0.6 to ~ -0.3, from the outermost to the innermost fields, using isochrones from VandenBerg et al. (2000), re-calibrated to match Galactic globular clusters; see Rejkuba et al. (2005) for more details. Note that the lack of high-metallicity stars in the 21-31-kpc field of NGC 5128 is a selection effect due to insufficiently deep V-band observations, so the shape of this MDF cannot be fairly compared with the inner and outer fields. Even so, in both galaxies, there are stars of at least

Solar metallicity, and there is evidence for super-metal-rich stars, particularly in the innermost field.

There are certain cautions to be kept in mind when considering CMD-inferred MDFs. First, old, single-age stellar populations are assumed in order to compute these distributions. This is due to the lack of age information in CMDs that are insufficiently deep: without CMDs that reach down below the main-sequence turn-off(s), the age is unknown (although the presence of bright AGB stars can indicate intermediate-aged populations). A shift, or worse, a distribution, in stellarpopulation age can change both the mean and the shape of the inferred MDF. This can be seen by comparing the MDFs of M32 inferred from the 10- and 15-Gyr-old isochrones in Figure 16.3 and imaging a mixture of these populations. Second, the assumed abundance pattern (as parametrised by, say, [a/Fe]) can also alter both the mean and the shape of the inferred MDF. This is because changes in the abundance pattern can alter the temperature of the stars on the giant branch (e.g. Salaris et al. 1993). On the other hand, this might not be a significant problem for these two galaxies: for M32, because its abundance pattern is very close to that of the Sun (Worthey 2004), and the isochrones used to infer its MDF have scaled-Solar composition (Worthey 1994); and for NGC 5128, because, although the galaxy probably has [a/Fe] > 0 because it is a giant elliptical (see below), the isochrones used to infer its MDFs were a-enriched (Rejkuba et al. 2005).

2.2 Abundances from planetary nebulae

Planetary nebulae (PNe) are the second-to-last phase of the life of all low- and intermediate-mass stars. Although this phase is quite short, the large flux of stars through it means that many PNe are visible at any given time. Because this phase is so short, their density is much lower than that of the field stars from which they evolved, and therefore they are much less crowded on the sky. They are emissionline objects and are fairly easy to distinguish from the absorption-line background of their hosts. Using standard emission-line techniques, it is possible to determine elemental abundances directly for each PNe.

However, PNe are quite faint. Special wide-field spectroscopy techniques have been developed to determine the kinematics of PNe systems around galaxies (e.g. Douglas et al. 2002), which utilise a narrow-band region around the [O iii]A5007 line. However, these techniques do not (yet) allow abundance determinations, since these require a larger wavelength region: oxygen abundances require measurements of the [O iii]A4363, 5007 and [S ii]A6716, 6731 lines. Moreover, the [O iii]A4363 line is very weak. Unfortunately, without this line, the electron temperature Te cannot be measured directly, and strong-line techniques (such as those used for H ii regions) cannot be used for PNe because the temperature of the central stars varies significantly from object to object, and photoionisation models do not give unique determinations (Stasinska et al. 2006). On the other hand, unlike in H11 regions, the [O iii]A4363 line is still visible for metal-rich PNe because of the high density of the nebulae and the high temperatures of the central stars (Stasinska et al. 2005). With large telescopes (>4 m) and sensitive spectrographs, bright PNe in early-type galaxies become reasonable tracer objects for the abundance patterns of their hosts.

Many early-type galaxies host well-known PNe systems, although have abundance measurements been attempted for only a handful of them. Here I will concentrate on the three cases that I know of for which independent abundance determinations have been obtained by other methods: M32 (Stasinska et al. 1998), NGC 5128 (Walsh et al. 1999, 2006) and NGC 4697 (Mendez et al. 2005). Of these, only in M32 and NGC 5128 have PNe with [O iii]A4363 lines been detected: nine PNe in M32 (Stasinska et al. 1998) and two (so far) in NGC 5128 (Walsh et al. 2006). None of the PNe in NGC 4697 show [O iii]A4363 individually, although it is detected in an average spectrum (Mendez et al. 2005). I will therefore (reluctantly) no longer consider it here, since the Te of the individual PNe cannot be determined and therefore their abundances are likely to be very uncertain. In M32, Stasinska etal. (1998) found ([O/H]) = -0.5, in reasonable agreement with, and at projected galactocentric distances similar to those of, the Grillmair et al. (1996) CMD study. Walsh et al. (1999) find a similar average [O/H] abundance ratio for NGC 5128 in a field close to the centre, although for none of these PNe has [O iii]A4363 been detected. Walsh et al. (2006) appear to confirm this result using two PNe with Te estimates and also find evidence for two distinct populations of PNe in NGC 5128, a 'bulge' and a 'halo' population.

Abundances from PNe appear to be a rich area for further study, but the results of Mendez et al. (2005) and Walsh et al. (2006) suggest that perhaps telescopes of the current generation are at their limits even for PNe in galaxies as close as NGC 5128 (3.8Mpc) (Rejkubaetal. 2005).

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