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Astron. Astrophys. 346, L1-L4 (1999)

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3. Atmospheric parameters

3.1. Fit procedure and model atmospheres

To derive effective temperatures, surface gravities and helium abundances we fitted the observed Balmer lines [FORMULA] to H10 (excluding [FORMULA] because of possible blending problems with the CaII H line) and the helium lines (HeI 4026, 4388, 4471, 4922 Å) with stellar model atmospheres. We corrected the spectra for radial velocity shifts, derived from the positions of the Balmer and helium lines and normalized the spectra by eye.

We computed model atmospheres using ATLAS9 (Kurucz 1991) and used Lemke's version of the LINFOR 1 program (developed originally by Holweger, Steffen, and Steenbock at Kiel University) to compute a grid of theoretical spectra which include the Balmer lines [FORMULA] to H22 and HeI lines. The grid covered the range 7,000 K [FORMULA] [FORMULA] [FORMULA] 35,000 K, 2.5 [FORMULA] [FORMULA] [FORMULA] 5.0, -3.0 [FORMULA] [FORMULA] [FORMULA] -1.0, at a metallicity of [M/H] = -1.5.

To establish the best fit we used the routines developed by Bergeron et al. (1992) and Saffer et al. (1994), which employ a [FORMULA] test. The fit program normalizes model spectra and observed spectra using the same points for the continuum definition. The results are plotted in Fig. 2 (upper panel). The errors are estimated to be about 10% in [FORMULA] and 0.15 dex in [FORMULA] (cf. Moehler et al. 1997). Representative error bars are shown in Fig. 2. To increase our data sample we reanalysed the NTT spectra described and analysed by Moehler et al. (1997). For a detailed comparison see Moehler et al. (1999a).

[FIGURE] Fig. 2. Temperatures and gravities of the programme stars in NGC 6752. upper panel : determined from models with cluster metallicity ([M/H] = -1.5), central panel : adopting a solar metallicity model stratification ([M/H] = 0) and spectrum synthesis with solar iron abundance but cluster abundances for all other metals M ([M/H]=-1.5) lower panel : adopting a super-solar metallicity model stratification ([M/H] = +0.5) and iron abundance ([Fe/H] = +0.5) but cluster abundances ([M/H] = -1.5) for all other metals in the spectrum synthesis. For more details see text. Also plotted are the zero-age HB (ZAHB) and terminal-age HB (TAHB, i.e., central helium exhaustion) from the Sweigart (1999) tracks for metallicity [M/H] = -1.56. The dashed and solid lines correspond to tracks with and without mixing, respectively. [FORMULA] measures the difference in hydrogen abundance X between the envelope ([FORMULA]) and the innermost point reached by the mixing currents ([FORMULA]) in the red giant precursors and is thus an indicator for the amount of helium mixed into the envelope of the red giant. Representative error bars are plotted

3.2. Iron abundances

Due to the spectral resolution and the weakness of the few observed lines a detailed abundance analysis (such as that of Behr et al., 1999) is beyond the scope of this paper. Nevertheless we can estimate the iron abundance in the stars by fitting the FeII lines marked in Fig. 1. A first check indicated that the iron abundance was about solar whereas the magnesium abundance was close to the mean cluster abundance.

As iron is very important for the temperature stratification of stellar atmospheres we tried to take the increased iron abundance into account: We used ATLAS9 to calculate a solar metallicity atmosphere. The emergent spectrum was then computed from the solar metallicity model stratification by reducing the abundances of all metals M (except iron) to the cluster abundances ([M/H] = -1.5). It was not possible to compute an emergent spectrum that was fully consistent with this iron-enriched composition, since the ATLAS9 code requires a scaled solar composition. We next repeated the fit to derive [FORMULA], [FORMULA], and [FORMULA] with these enriched model atmospheres. The results are plotted in Fig. 2 (central panel).

For each star observed at the ESO 1.52m telescope we then computed an "iron-enriched" model spectrum with [FORMULA], [FORMULA] as derived from the fits of the Balmer and helium lines with the "enriched" model atmospheres (cf. Fig. 2, central panel) and [FORMULA] = -2. The fit of the iron lines was started with a solar iron abundance and the iron abundance was varied until [FORMULA] achieved a minimum. As the radiative levitation in BHB stars is due to diffusion processes (which is also indicated by the helium deficiency found in these stars) the atmospheres have to be very stable. We therefore kept the microturbulent velocity [FORMULA] at 0 km s-1 - the iron abundances plotted in Fig. 3 are thus upper limits. The mean iron abundance turns out to be [Fe/H] [FORMULA] dex (for 18 stars hotter than about 11,500 K) and [FORMULA]1.6 for the one star cooler than 11,500 K. Although the iron abundance for the hotter BHB stars is about a factor of 50 larger than the cluster abundance, it is smaller by a factor of 3 than the value of [Fe/H] = +0.5 estimated by Grundahl et al. (1999) as being necessary to explain the Strömgren u-jump observed in u, [FORMULA] colour-magnitude diagrams.

[FIGURE] Fig. 3. The iron and helium abundances for the stars observed with ESO 1.52m telescope. Iron was not detected in the coolest star and is plotted as an upper limit. The trend to lower helium abundances for higher temperatures agrees with the findings of Behr et al. (1999). Iron is obviously enhanced to roughly solar abundances. The mean iron abundance as derived from our spectra ([Fe/H] = +0.13) and the cluster abundance ([Fe/H] = -1.54) are marked. The asterisk marks the results of Glaspey et al. (1989) for the hotter of their two BHB stars in NGC 6752

Our results are in good agreement with the findings of Behr et al. (1999) for BHB stars in M 13 and Glaspey et al. (1989) for two BHB stars in NGC 6752. Again in agreement with Behr et al. (1999) we see a decrease in helium abundance with increasing temperature, whereas the iron abundance stays roughly constant over the observed temperature range.

3.3. Influence of iron enrichment

From Fig. 2 it is clear that the use of enriched model atmospheres moves most stars closer to the zero-age horizontal branch (ZAHB). The three stars between 10,000 K and 12,000 K, however, fall below the canonical ZAHB when fitted with enriched model atmospheres. This is plausible as the radiative levitation is supposed to start around 11,500 K (Grundahl et al. 1999) and the cooler stars therefore should have metal-poor atmospheres (see also Fig. 3 where the coolest analysed star shows no evidence of iron enrichment). We repeated the experiment by increasing the iron abundance to [Fe/H]=+0.5 (see Fig. 2 lower panel), which did not change the resulting values for [FORMULA] and [FORMULA] significantly.

Since HB stars at these temperatures spend most of their lifetime close to the ZAHB, one would expect the majority of the stars to scatter (within the observational error limits) around the ZAHB line in the [FORMULA], [FORMULA]-diagram. However, this is not the case for the canonical ZAHB (solid lines in Fig. 2) even with the use of iron-enriched model atmospheres (central and lower panels in Fig. 2). The scatter instead seems more consistent with the ZAHB for moderate helium mixing (dashed lines in Fig. 2). Thus the physical parameters of HB stars hotter than [FORMULA] K in NGC 6752, as derived in this paper, are best explained by a combination of helium mixing and radiative levitation effects.

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© European Southern Observatory (ESO) 1999

Online publication: May 6, 1999
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