M. Zoccali1,2 - B. Barbuy3 - V. Hill4 - S. Ortolani5 - A. Renzini6 - E. Bica7 - Y. Momany5 - L. Pasquini6 - D. Minniti1 - R. M. Rich8
1 - Universidad Catolica de Chile, Department of Astronomy & Astrophysics,
Casilla 306, Santiago 22, Chile
2 -
Princeton University Observatory, Peyton Hall, Princeton, NJ 08544, USA
3 -
Universidade de São Paulo, IAG, Rua do Matão 1226,
Cidade Universitária, São Paulo 05508-900, Brazil
4 -
Observatoire de Paris-Meudon, 92195 Meudon Cedex, France
5 -
Università di Padova, Dipartimento di Astronomia, Vicolo
dell'Osservatorio 2, 35122 Padova, Italy
6 -
European Southern Observatory, Karl Schwarzschild Strasse 2, 85748
Garching bei München, Germany
7 -
Universidade Federal do Rio Grande do Sul, Departamento de Astronomia,
CP 15051, Porto Alegre 91501-970, Brazil
8 -
UCLA, Department of Physics & Astronomy, 8979 Math-Sciences Building,
Los Angeles, CA 90095-1562, USA
Received 2 April 2004 / Accepted 10 May 2004
Abstract
High resolution spectra of five stars in the bulge globular cluster
NGC 6528 were obtained at the 8m VLT UT2-Kueyen telescope with the
UVES spectrograph. Out of the five stars, two of them showed evidence
of binarity. The target stars belong to the horizontal and red giant
branch stages, at
K. Multiband V, I, J, H,
photometry was used to derive initial effective
temperatures and gravities. The main purpose of this study is the
determination of metallicity and elemental ratios for this template
bulge cluster, as a basis for the fundamental calibration of
metal-rich populations. The present analysis provides a metallicity
and the
-elements O, Mg and Si, show
[
/Fe]
+0.1, whereas Ca and Ti are around the solar
value or below, resulting in an overall metallicity Z
.
Key words: Galaxy: bulge - globular clusters: individual: NGC 6528 - stars: abundances - stars: atmospheres
The globular cluster NGC 6528 is located in Baade's Window, at (J2000)
,
(
,
), at a distance
kpc from the Sun and
kpc from the
Galactic center (Barbuy et al. 1998).
Ortolani et al. (1992) first presented BVRI CCD Colour-Magnitude Diagrams (CMD) of NGC 6528. Ortolani et al. (1995) showed that NGC 6528, together with its "twin'' NGC 6553, has a Colour Magnitude Diagram (CMD) and a luminosity function very similar to that of the Galactic bulge, as seen in the Baade's Window field. This is a clear indication that the two populations have comparable age and metallicity.
For this reason, the two bulge globular clusters (GCs) NGC 6528 and NGC 6553 have often been used as templates for the old, metal-rich population of the Milky Way central spheroid (e.g. Zoccali et al. 2003).
Despite its key importance for the interpretation of the formation of
our own Galaxy and, by extension, of external spheroids, the absolute
value of the Fe and -element abundance, for both NGC 6528 and
NGC 6553 still are a matter of debate. The main reasons for this are:
i) the intrinsic difficulty to observe these faint stars; ii) the ambiguous location of the continuum, due to the severe line
crowding in these metal-rich objects, iii) the presence of
strong molecular (mainly TiO) bands in the brightest, coolest stars,
together with effects of
-element enhancements on the
continuum absorption; iv) the uncertainty of the effective
temperature scale for stars as hot as the Horizontal Branch (HB):
temperatures derived by imposing excitation equilibrium for the
Fe I lines tend to be overestimated probably due to blends, and
possibly to NLTE effects on Fe I in giants, whereas
temperatures derived from colours show uncertainties due to reddening
variations.
A few stars in each of the two GCs have been observed in the past few years at high spectral resolution, but derived metallicities show discrepancies, higher than the estimated uncertainty, thus revealing that there are systematic errors not properly taken into account.
In NGC 6553, two were observed with CASPEC, at the ESO 3.6 m Telescope,
by Barbuy et al. (1999). They derived
,
,
and
.
A significantly higher metallicity was
obtained instead by Cohen et al. (1999), from HIRES@Keck I
observations of 5 HB stars:
,
.
More recently, Origlia et al. (2002)
measured the metallicity of NGC 6553 from near-IR spectra obtained
with NIRSPEC@Keck II, obtaining
,
with
.
Meléndez et al. (2003) used PHOENIX@Gemini-South
to obtain spectra of
R = 50 000 in the H band, and they derived
and
.
A summary of the previous metallicity
determinations of NGC 6553 is given in Table 5 of Barbuy et al. (1999).
For NGC 6528 the only high-resolution analysis published so far has
been that of Carretta et al. (2001, hereafter C01), who observed four
HB stars, with HIRES at Keck I. They obtained
and
.
Recently, Momany et al. (2003) determined the metallicity of NGC 6528
from RGB morphology indicators, in a (K, V-K) CMD, obtaining values
in the range
depending on the adopted
calibration.
In this work we present detailed abundance analysis of three stars in
NGC 6528 using high resolution échelle spectra obtained with UVES at
the ESO VLT-UT2 Kueyen 8 m telescope, in an attempt to reduce the
discrepancy between previous measurements using higher spectral
resolution (
)
relative to previous studies (e.g.
in C01 and
in Barbuy et al. 1999). The
analysis is restricted to relatively hot (
K)
stars, in order to avoid the above mentioned problems with molecular
bands.
The paper is organized as follows. Spectroscopic observations are described in Sect. 2. Section 3 focuses on the determination of the stellar parameters: effective temperature and gravity from photometry and spectroscopy. Equivalent width and oscillator strenghts are discussed in Sect. 3. Iron abundance is derived in Sect. 4 and abundance ratios in Sect. 5. The results are then discussed in Sect. 6.
V and I photometry of the central region of NGC 6528 was obtained
using two sets of WFPC2 observations with the Hubble Space Telescope
(HST), (GO5436: Ortolani et al. 1995; GO8696: Feltzing &
Johnson 2002). Infrared
observations were obtained with
the SOFI infrared camera of the ESO New Technology Telescope
(NTT). The details about both optical and near-IR data are given in
Momany et al. (2003).
Five isolated stars were selected from our sample for spectroscopic observations (Fig. 1). The target location on both the SOFI near-IR and the HST optical CMD is shown in Fig. 2.
Table 1 lists the coordinates, magnitudes and colours of the sample stars. The star's identifications follow the notation by van den Bergh & Younger (1979), where the prefix I corresponds to the internal ring. For each star the first line corresponds to the original magnitudes, while the second line corresponds to the magnitudes corrected for total (see Sect. 3.1) and differential extinction. The latter correction has been carried out following the method described in Piotto et al. (1999) and Zoccali et al. (2001). This method assumes that the shift relative to the RGB fiducial mean locus, between different subregions, is dominated by differential reddening. The observed field is divided into small subregions, and the reddening of each of them is estimated by comparison with the fiducial mean locus of NGC 6528 from Momany et al. (2003). The result of this procedure for the (V, V-I) CMD is shown in Fig. 3.
![]() |
Figure 2: The SOFI near-IR ( left) and HST optical ( right) CMD of NGC 6528 with the target stars shown as big filled circles. |
A posteriori, we identified some of our targets also in the field of
view of a J,H set of NICMOS@HST data (Ortolani et al. 2001). The
spatial resolution of the NICMOS observations allowed us to discover
that I-16 has in fact a nearby companion which contaminates the
spectrum by about 23
in the J and H bands (see Fig. 4).
For this reason,
I-16 was discarded from the spectroscopic analysis.
High resolution spectra for five members of NGC 6528, in the
wavelength range
4800-6800 Å, were obtained
with the UVES spectrograph at the ESO VLT. The target membership had
been previously verified from radial velocities by Coelho et al. (2001). In what follows we will consider mainly the reddest
portion of the spectrum (5800-6800 Å) covered by the MIT
backside illuminated and AR coated CCD ESO # 20 of
pixels,
of pixel size
m. The bluer region was used only to derive
the carbon abundance from the C2 (0, 1) bandhead at
5635 Å. With the UVES standard setup 580, the resolution
is
for a 1 arcsec slit width, and
for a
slit of 0.8 arcsec. The pixel scale is 0.0174 Å/px.
The log of the observations is shown in Table 2.
The spectra were reduced using the UVES context of the MIDAS reduction package, including bias and inter-order background subtraction, flatfield correction, extraction and wavelength calibration (Ballester et al. 2000). When possible we used the optimal extraction routine, rejecting the cosmic rays. However, in a few cases (e.g. star I-16 and I-18, which were observed through the same slit), the optimal extraction could not be used, and therefore we turned to the standard extraction routine. The two spectra obtained for each target star were summed before the analysis.
A mean
km s
or heliocentric
km s
was found for NGC 6528,
higher than the values of 160 km s
given in Zinn & West
(1984) or 143 km s
given in Zinn (1985), but in very good
agreement with more recent values of
km s
and
km s
measured respectively by Minniti (1995)
and Rutledge et al. (1997).
A reddening of
E(B-V) = 0.55 was given in Momany et al. (2003),
however this was a mean value from determinations relative to NGC 6553
and 47 Tucanae, in different colours (see their Table 3), and those
values are somewhat dependent on the metallicity and -element
assumptions in isochrone models. They can be used as guidelines, but
for temperature determinations we need to investigate further about
the value to adopt, since a
implies in a
K.
Reddening vs. spectral type:
A very well established result, after the analysis of Schmidt-Kaler
(1961), is that the absorption (in the UBV bands) is a function of
the energy distribution of the star within the broadband filters (see
also more recent discussions and calculations in Olson 1975; Grebel
& Roberts 1995; Fitzpatrick 1999). Such effect is due to the
decrease of the extinction, toward longer wavelengths, across the width
of the filter. The net effect results in a shift of the effective
wavelength, for a reddened star, to a longer wavelength than for an
unreddened star. The magnitude of the shift depends on the energy
distribution (therefore on the spectral type) of the star. This makes
redder (cooler) stars systematically less absorbed. The effect is
more pronounced in the B band than in V, for the same amount of
interstellar dust, and thus the colour excess becomes sensitive to the
stellar energy distribution. Schmidt-Kaler (1961) introduced a
parameter
to quantify this effect:
NGC 6522 is another globular cluster located in the Baade's Window
that can be used to derive colour excess in the region. The work by
Terndrup & Walker (1994) on several fields in NGC 6522 gives
E(B-V)=0.52 for B-V=0 stars, corresponding to about 0.44-0.45 for our
K-M stars. This is done via a simple fit to CMDs relative to other
globular clusters having the same RGB morphology. Terndrup et al.
(1998) derived Av = 1.4 on a proper motion cleaned CMD of NGC 6522. This
gives again about
E(B-V)=0.45 for K-M stars. We can conclude that
there is quite strong evidence that NGC 6522 has a reddening close to
0.44-0.45. Now we should estimate the difference between NGC 6528 and
NGC 6522. A check on Stanek's (1996) maps indicates that the two
regions have very similar extinctions. The zero point has been
revised by Alcock et al. (1998) and Gould et al. (1998). Following
Alcock et al., from their Table 1, the extinction near NGC 6528 (at
04
29.46
,
,
J2000.0) is
mag higher than near NGC 6522. Using a total-to-selective absorption
,
this transforms into
E(B-V)= 0.01- 0.02mag, which makes a very small difference. The same authors derived
E(B-V)=0.48 for (B-V)=0.0 stars for the less redenned side of the
Baade Window, corresponding to one of the lowest values (about 0.42
for K-M stars). They strongly argue against high values. Since they
start from luminosity measurements of the RR Lyrae stars, they have to
make assumptions about the total-to-selective absorption
/E(B-V) value to get the E(B-V), but still this appears to
be a quite firm result.
Comparison cluster | [Fe/H] |
![]() |
M 13 | -1.54 | 0.60 |
NGC 6229 | -1.43 | 0.50 |
NGC 6723 | -1.12 | 0.50 |
NGC 6681 | -1.51 | 0.60 |
NGC 6266 | -1.29 | 0.42 |
NGC 5904 | -1.27 | 0.50 |
NGC 1851 | -1.22 | 0.47 |
NGC 1904 | -1.57 | 0.55 |
We carried out an independent reddening measurement for NGC 6522 using
HST photometry from Piotto et al. (2002). The CMD of NGC 6522 was
compared to those of reference clusters, having comparable
metallicities, and which have been observed with the same HST
instrumentation. The reddening difference between NGC 6522 and each
reference cluster is measured as the colour difference of their red
giant branches,
,
as shown in
Table 3.
A cluster metallicity can be estimated from the curvature of the RGB,
since more metal-rich giants show progressively fainter RGB magnitudes
(as well as redder RGB colours), due to higher blanketing (Ortolani et al. 1990, 1991). The metallicities reported are from Harris (1996, as
updated in
http://physun.physics.mcmaster.ca/Globular.html).
Note that the more metal rich clusters will give obviously a bit lower
reddening and viceversa: M 13, NGC 6681 and NGC 1904 appear to have a
more metal poor morphology as concerning the RGB, and give higher
.
The
would be given by comparing
two clusters of same metallicity, given that there should be no
blanketing differences between them. The best morphological fits are
obtained for NGC 5904 (M 5), NGC 1851 and NGC 6266. Harris (1996) gives
for NGC 6522, and in these cases the derived
is close to the E(B-V) value. In conclusion, it seems
that
E(B-V)=0.42-0.50 represents a fiducial reddening value for NGC
6522 for the colour of horizontal branch to G-K stars.
In order to obtain the reddening for NGC 6528 we would have to add the
reddening difference between the former and NGC 6522 (+0.02), but
another
would be needed to translate this for
K-M stars (cf. parameter
defined above). Therefore, for our
stars, we adopt the same reddening
determined
above for NGC 6522. Note that this result is independent of possible
zero point errors in the Piotto et al. photometry because it is based
on differential photometry.
Yet another free parameter comes into play, because the extinction law
towards the bulge is somewhat controversial. Extremely low values of
the total-to-selective absorption
were recently proposed by e.g. Udalski (2003) in the galactic bulge
OGLE II field, as opposed to the traditional
from
Cardelli et al. (1989) or Rieke & Lebofsky (1985). For completeness,
let us recall that the total-to-selective absorption
depends on the amount of reddening (Olson 1975) and on metallicity
(Grebel & Roberts 1995), as discussed in Barbuy et al. (1998), and
the value should be somewhat higher for high reddening and high
metallicity values. If lower
values are assumed, our
resulting dereddened colours will be redder, and therefore the deduced
effective temperatures will be cooler.
For the present analysis we adopt
E(B-V) = 0.46 for NGC 6528 and the
extinction law given by Dean et al. (1978) and Rieke & Lebofsky
(1985), namely:
,
E(V-I)/E(B-V)=1.33,
E(V-K)/E(B-V)=2.744 and
E(J-K)/E(B-V)=0.527 (Table 1).
AAM99 | |||||||
Star |
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log g |
I-18 | 4504 | 4520 | 4409 | 4511 | -0.47 | -0.16 | 2.0 |
I-36 | 4071 | 4051 | 3937 | 4019 | -0.87 | -0.93 | 1.5 |
I-42 | 3897 | 3975 | 3937 | 3936 | -0.97 | -0.93 | 1.5 |
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Figure 5: Calibrations of Alonso et al. (1999), of V, I, J, H and Kas a function of effective temperature (solid lines), where the magnitudes for the sample stars are overplotted (filled squares). |
Effective temperatures were derived from V-K, V-I and J-K using
the colour-temperature calibrations by Alonso et al. (1999) together
with the reddening laws by Dean et al. (1978) and Rieke & Lebofsky
(1985). The
derived from each colour are listed in
Table 4. We have excluded from the spectroscopic analysis
the star I-16 as explained in Sect. 2.1, and also star I-24, being a
suspect unresolved binary, for the reasons explained below (Sect. 4).
In order to verify that all the colours are compatible with the
adopted temperature, we compared the spectral energy distribution
(SED) of our targets with the prediction by Alonso et al. (1999).
Figure 5 shows the flux obtained from the V,I,J,H and Kmagnitudes, together with Alonso et al.'s calibrations for stars of
different temperatures (4000-5250 K). Since Alonso et al. give a
relation between colours and temperatures, the predicted flux in each
band was obtained by assuming an arbitrary value for the Vmagnitude. A vertical shift was then allowed to match the SEDs of the
sample stars.
The classical relation
)
was used, adopting
K,
and
(Cram 1999).
A distance modulus of
(m-M)0 = 15.10 was adopted from Harris (1996)
together with a reddening of
E(B-V) = 0.46 and
,
as
discussed in the previous sections. The bolometric corrections from
Alonso et al. (1999) and corresponding gravities are given in
Table 4.
Spectroscopic gravities were also derived by imposing ionization equilibrium between FeI and FeII lines (see Sect. 4).
The equivalent widths have been measured using a new automatic code developed by P. Stetson (DAOSPEC, Stetson et al. 2004, in preparation). This routine fits a Gaussian profile to all lines, imposing a constant FWHM. In order to take into account the broader profile of strong lines, we manually measured lines stronger than 140 mÅ allowing a best fit for the FWHM of each line. Star I-24 shows lines with FWHM significantly larger than that of the other stars (a factor of 1.5), and most of the lines have a systematic asymmetry (Fig. 6). This was interpreted as evidence of the presence of an unresolved companion, and for this reason the star was excluded from the following analysis. Note that at lower spectral resolution such unresolved binaries would not have been detected, and therefore the results would not have reflected the true stellar abundances.
The Fe I line list and respective oscillator strengths given in NIST (Martin et al. 2002) and some lines with gf-values from McWilliam & Rich (1994, hereafter MR94) were used to derive spectroscopic parameters. Eight measurable Fe II lines, and their respective oscillator strengths from Biémont et al. (1991), and renormalized by Meléndez & Barbuy (2004, in preparation), were used to obtain ionization equilibrium.
The damping constants were computed where possible, and in particular
for most of the Fe I lines, using the collisional broadening
theory of Barklem et al. (1998, 2000, and references therein). For
lines of Si, Ca and Ti the derived values of the interaction constant
C6, where C6 = 6.46E-34 (/63.65)5/2,
being
the cross section of collision computed by Barklem et al., which is
related to the collisional damping constant
=
17v3/5C62/5N, are reported in Table 6.
Otherwise a fit to the solar spectrum or standard van der Waals values
(Unsöld 1955; Gray 1976) are employed.
The oscillator strengths of Na, Mg, Si, Ca and Ti lines given by Brown
& Wallerstein (1992), MR94, the NIST database (Martin et al. 2002),
Barbuy et al. (1999) from fits to the solar spectrum, the VALD
database (Kupka et al. 1999), and the recent compilation by Bensby et al. (2003) are reported in Table 6. The adopted
oscillator strengths for these lines were then selected from a fit to
the solar spectrum observed with the same VLT-UVES instrumentation,
adopting a NMARCS solar atmospheric model (Edvardsson et al. 1993),
in order to be compatible with the model grid used to analyse the
sample giants.
These gf-values were checked against the spectrum of
Arcturus (Hinkle et al. 2000), by adopting the stellar parameters
derived by Meléndez et al. (2003), namely
K,
log g = 1.55,
,
km s-1 and
abundances
,
,
,
and other
abundance ratios from the literature, and rederived from our fits:
,
,
,
,
,
.
The selected gf-values are marked in
bold face in Table 6.
For the oxygen forbidden line [OI]6300 Å we adopt the oscillator strength recently derived by Allende Prieto et al. (2001): log gf = -9.716.
For lines of the heavy elements BaII, LaII and EuII, a hyperfine structure was taken into account, based on the hyperfine constants by Lawler et al. (2001a) for EuII, Lawler et al. (2001b) for LaII and Biehl (1976) for BaII.
Solar abundances were adopted from Grevesse et al. (1996), except for
the value for oxygen where (O) = 8.77 was assumed, as
recommended by Allende Prieto et al. (2001) for the use of 1D model
atmospheres.
Photospheric 1D models for the sample giants were extracted from the NMARCS grid (Plez et al. 1992), originally developed by Bell et al. (1976) and Gustafsson et al. (1975).
The LTE abundance analysis and the spectrum synthesis calculations were performed using the codes by Spite (1967, and subsequent improvements made in the last thirty years), Cayrel et al. (1991) and Barbuy et al. (2003).
The stellar parameters were derived by initially adopting the
photometric effective temperature and gravity, and then further
constraining the temperature by imposing excitation equilibrium for
Fe I lines and the gravity by imposing ionization equilibrium
for Fe I and Fe II. Microturbulence velocity
was
determined by canceling the trend of Fe abundance vs. equivalent
width.
The same method and the same line list, applied to the Sun and
Arcturus, yield metallicities of [Fe I/H] = 0.09,
[Fe II/H] = 0.01 and [Fe I/H] = -0.49, [Fe II/H] = -0.56 respectively. Note that in this case the stellar parameters
that we converged to are
K, log g = 4.4,
km s-1 for the Sun and
K, log g =
1.8,
km s-1 for Arcturus, in good agreement
with expectations.
Table 5 reports [Fe I/H] and [Fe II/H] values
for the sample stars. Figures 7-9 illustrate the results of the analysis for stars I-18,
I-36 and I-42 respectively. The abundances from Fe II lines
for star I-42 showed relatively large scatter, possibly due to a
combination of the lower S/N and cooler temperature. We therefore
compared the equivalent width of Fe II lines in I-42 with those
of I-36, having very similar
,
and rejected the lines
showing large deviations from the mean percent difference. This
allowed us to have a smaller dispersion in [Fe II/H] and
therefore to get a spectroscopic constraint on the gravity for I-42.
A second set of parameters based on the photometry was derived by imposing the temperature and gravity values from Table 4 and determining the metallicity and the microturbulence velocity from Fe I and Fe II. The two sets of stellar parameters are given in Table 5.
As a check ionization equilibrium was also used as a temperature indicator, as suggested by McWilliam & Rich (2003, hereafter MR03), keeping the gravity fixed to its photometric value, and varying the temperature to recover the Fe ionization equilibrium. This leads to temperatures in better agreement with the photometric values, while keeping excitation equilibrium virtually identical to the formally best fitting spectroscopic temperature.
This shows that the true temperature of the sample stars is very
likely hotter (by 200 K) than the photometric one. In the
following we will keep this last set of parameters as the one which
gives the most self consistent picture. The errors on [Fe/H] were
estimated as follows. The formal error on the slope in the FeI vs.
ionization potential implies an error in the temperature of
100 K
for all the stars. From the abundance analysis of Arcturus, we
estimated that a variation of
100 K in the temperature
propagates into a difference of
0.07 in [FeI/H] and
0.06 in
[FeII/H]. This variation in the metallicity of Arcturus was found by
forcing excitation equilibrium (i.e, by fixing the gravity) and canceling
the trend in abundance with EW (i.e., by fixing
)
with the
new (uncertain by
100 K) temperatures. This procedure gives an
estimate of the systematic error due to a possible uncertain
assumption on the input parameters. The statistical error, measured
from the spread in abundance obtained from different lines, is
for the three stars. The final error,
being the quadratic sum of the systematic and statistical components,
is quoted in Table 5.
Star |
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Spectroscopic: | |||||
I-18 | 4800 | 2.3 | +0.05 (![]() |
-0.07 (![]() |
1.5 |
I-36 | 4300 | 1.8 | -0.04 (![]() |
-0.08 (![]() |
1.5 |
I-42 | 4200 | 1.6 | -0.16 (![]() |
- | 1.3 |
Photometric: | |||||
I-18 | 4500 | 2.0 | -0.05 (![]() |
+0.18 (![]() |
1.3 |
I-36 | 4000 | 1.5 | +0.00 (![]() |
+0.35 (![]() |
1.5 |
I-42 | 4000 | 1.5 | -0.19 (![]() |
- | 1.3 |
Ionization temperature: | |||||
I-18 | 4700 | 2.0 | -0.05 (![]() |
-0.11 (![]() |
1.5 |
I-36 | 4200 | 1.5 | -0.13 (![]() |
-0.09 (![]() |
1.5 |
I-42 | 4100 | 1.6 | -0.14 (![]() |
-0.08 (![]() |
1.2 |
![]() |
Figure 8: Same as Fig. 7 for star I-36. |
![]() |
Figure 9: Same as Fig. 7 for star I-42. |
Elemental abundances were obtained through line-by-line spectrum synthesis calculations, for the line list given in Tables 6 and 7.
The calculations of synthetic spectra were carried out as described in
Cayrel et al. (1991) and Barbuy et al. (2003), where molecular lines
of CN (A2-X2
), C2 Swan (A3
-X3
), TiO (A3
-X3
)
and TiO (B3
-X3
)
' systems are taken into account. As a first guess, the
abundances obtained from the equivalent width, in the "classical''
way, without taking into account the effect of molecules, were
adopted. These first guess values are listed in Table 8.
Species |
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log gf | I-18 | I-36 | I-42 |
[OI] | 6300.311 | 0.00 | -9.716 | +0.10 | +0.15 | +0.20 |
ScII | 6300.67 | 1.50 | hfs | -0.25 | -0.20 | +0.00 |
C(C2) | 5634.2, 35.5 | (0, 1) | - | 0.00 | 0.00 | 0.00 |
N(CN) | 6332.2 | (5, 1) | - | +0.60 | +0.60 | +0.60 |
BaII | 6141.727 | 0.70 | -0.08 | -0.05 | +0.20 | -0.05 |
BaII | 6496.908 | 0.60 | -0.38 | 0.00 | +0.10 | 0.00 |
LaII | 6390.480 | 0.32 | -1.45 | -0.30 | -0.15 | +0.00 |
EuII | 6645.127 | 1.38 | +0.12 | +0.15 | +0.15 | +0.15 |
The comparison between the synthetic spectrum, generated including the
effects of molecules, and the observed one, for each line in
Table 6, allowed us to set a more realistic estimate of
the true abundance ratios. Our best values are listed in
Table 9. The comparison between Tables 8 and 9 gives an estimate of the importance of the effect
of molecule formation in the atmosphere of each star. The main effect
is a lowering of the continuum as illustrated in Fig. 10, where
TiO was computed for
and +0.2 (the latter value
is the adopted one), and for
,
4100 and 4200 K.
Errors in the abundance ratios can be due essentially to gf-values and damping constants used. They can be deduced from Table 6, where a difference of 0.2 dex in log gfessentially means a difference of 0.2 dex in abundance, for a given damping constant.
The carbon abundance
was derived for the 3 sample stars,
based on the C2(0, 1) 5635 Å, whereas the nitrogen abundance
was derived from a series of CN bandheads along the
spectrum. The region of the C2 line is very noisy, and the whole
region is blended with CN lines as well, such that the carbon
abundance is not as precise as all the other derivations in this work.
On the other hand, C+N is reliable. The N overabundance corresponds to
a mixing that is expected in giants, bringing carbon down and nitrogen
up, but since C+N is high, it does possibly involve as well a
primordial enhancement in C+N elements, as detected before in NGC 6553
(Meléndez et al. 2003).
Species | No. | I-18 | I-36 | I-42 |
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2 | -0.15 | +0.00 | +0.15 |
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2 | +0.62 | +0.30 | +0.18 |
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4 | +0.24 | +0.15 | +0.31 |
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3 | +0.12 | +0.08 | +0.11 |
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7 | -0.30 | -0.41 | -0.28 |
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6 | -0.30 | -0.08 | +0.23 |
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1 | -0.15 | - | +0.26 |
Species |
No. | I-18 | I-36 | I-42 | mean | C01 | MR94 |
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50 | -0.05 | -0.13 | -0.14 | -0.11 | +0.07 | - |
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1 | +0.10 | +0.15 | +0.20 | +0.15 | +0.07 | +0.03 |
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1 | +0.35 | +0.35 | +0.20 | +0.30 | +0.17 | - |
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2 | +0.60 | +0.40 | +0.30 | +0.43 | +0.40 | +0.20 |
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4 | +0.05 | +0.05 | +0.10 | +0.07 | +0.14 | +0.35 |
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3 | +0.05 | +0.10 | +0.10 | +0.08 | +0.36 | +0.14 |
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7 | -0.40 | -0.40 | -0.40 | -0.40 | +0.23 | +0.14 |
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6 | -0.30 | -0.20 | +0.20 | -0.10: | +0.03 | +0.37 |
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1 | -0.20 | -0.20 | +0.05 | -0.12: | - | - |
Figures 11-13 show examples of the
fit of oxygen, calcium and titanium lines, respectively, in the three
sample stars. While the abundance of [O/Fe] and [Ca/Fe] is very
homogeneuos in the three stars, Titanium is considerably higher in
star I-42. Figure 13 shows, for comparison, also a model
with low Ti (
)
with the stellar parameters of
I-42. Clearly such low Ti for I-42 can be excluded from the present
data.
The fit of the oxygen line [OI]6300.311 Å shown in
Fig. 11 also shows the nearby ScII 6300.67 Å line. Sneden &
Primas (2001) pointed out that the ratio of [OI] to ScII lines is a
more reliable measure of the oxygen enhancement. Indeed, OI and ScII
are the dominant species of these elements in cool stellar
atmospheres, and excitation potentials of the [OI]6300.31 Å
(
eV) and that of ScII 6300.67 Å (
eV) are similar. We used a hyperfine structure splitting and
corresponding log gf of the ScII line as given in Spite et al.
(1989). The present analysis gives
,
+0.35 and +0.20
for I-18, I-36 and I-42, respectively.
The line-by-line abundance ratios are given in Table 6, while the final resulting abundance ratios for the 3 sample stars, together with their mean, are shown in Table 9. Column 7 lists the results for the 4 HB stars analysed by C01 and Col. 8 shows the average of Baade's Window 11 K giants analysed by MR94. There is a general good agreement with respect to C01, except for a difference of 0.2 dex in [FeI/H], and the Si, Ca and BaII abundances which are higher in their study. The Mg, Al and Ti abundances found by MR94 in Baade's Window are higher than the values found in the present paper and in C01.
We have derived element abundances for three giants in the globular cluster NGC 6528. Five targets were originally selected. However, the high spatial resolution of the HST/NICMOS imaging frames, and the high UVES spectral resolution showed that two of them are in fact a blend of two stars (I-24 is likely a physical binary), and for this reason the two objects have been excluded from the analysis.
This case shows how strong can be the effect of projection or physical binaries in similar studies, when high spatial resolution images of the field are not available.
From the abundance analysis we conclude that Iron is about solar:
.
The odd-Z element sodium,
built up during carbon burning, is strongly overabundant
(
).
The -elements show a quite peculiar behaviour.
O, Mg and Si are just slightly above the solar value
(
)
while Ca and Ti are
underabundant with respect to the Sun (
). It is also somehow puzzling that despite the very
good agreement between the abundances derived from different lines in
the same star (Table 6), the three stars show a
spread in [Ti/Fe]. In particular, star I-42 shows a
higher Ti abundance with respect to the other two. Given
its radial velocity, the similarity with I-18 and I-36 in the
abundances of all the other elements, and the agreement between its
photometric and spectroscopic temperature, we tend to discard the
hypothesis of it being a field star just accidentally projected on the
cluster RGB sequence. We rather believe that its low temperature makes
stronger the effect of the TiO molecules in its atmosphere, which also
increases our uncertainty in the abundance analysis.
The comparison between our values and the ones in Carretta et al. (2001) reveals some discrepancies in Si, Ca, and Ti. While the
different treatment of molecules in the two analyses (Carretta et al.
did not account for them) might explain part of the discrepancy, the
large (0.6 dex) difference in the Ca abundance has to arise from some
other factor. From an independent analysis of the equivalenth widths
of the NGC 6528 stars observed by Carretta et al. (2001), MR03
found
,
much closer to the value found
here. Comparing the three analyses, it seems that the most probable
source of discrepancy are the adopted broadening factors. Indeed, it
must be noted that all the CaI lines are rather strong, in the saturation
regime of the curve of growth.
In fact the abundances of Ca and Ti do not always follow the other
-elements, as is the case for example with 47 Tucanae
where Ca-to-Fe is solar, and the other
s are overabundant
(Brown & Wallerstein 1992; Alves-Brito et al. 2004). According to
Woosley (private communication), there is metallicity sensitivity
in the yields because the mass loss varies with metallicity.
The yields therefore strongly depend on the mass loss, such that
Ca and Ti show variable proportions in different sites (Woosley
& Weaver 1995), and [Ti/Si] is often low.
From our analysis we find a modest enhancement of O, Mg and Si, which
is consistent with the trend of a lower -enhancement for the
metal-rich end of the bulge field stars found by McWilliam & Rich
(2003). However, the predictions by Matteucci & Brocato (1990) and
Matteucci et al. (1999) give higher expected values for O
and
at [Fe/H] = 0.0, and
somewhat less for Si and Ca (
). On
the other hand, the low values of
might indicate that there has been a deficiency
of SNII of low mass, as pointed out in MR03,
and according to predictions by Woosley & Weaver (1995).
It is also worth noticing
that Ca underabundance is very common in elliptical galaxies
(e.g. Thomas et al. 2003 for a discussion of possible causes).
The s-process elements Ba, La tend to show underabundances which is expected for an old population. The s-processing takes place in massive stars (Raiteri et al. 1993), as well as in relatively short lived and long lived intermediate mass stars (Gallino et al. 1998), indicating that there has been no significant contribution from long lived stars.
Acknowledgements
We thank the anonymous referee for his comments, that improved the readability of the manuscript. We are grateful to Stan Woosley for helpful comments on nucleosynthesis and chemical evolution of
elements, and to Cristina Chiappini for comments on chemical evolution predictions. B.B. acknowledges grants from CNPq and Fapesp. D.M. acknowledges grants from FONDAP Center for Astrophysics 15010003.
Star |
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V | I | J | H | ![]() |
||||
E(B-V) | V0 | I0 | J0 | H0 |
![]() |
![]() |
![]() |
(V-K)0 | (J-K)0 | |||
I-16 | 18:04:49.5 | -30:03:03.5 | -0.039 | 16.73 | 14.94 | 13.41 | 12.57 | 12.40 | ||||
0.421 | 15.42 | 14.19 | 13.12 | 12.34 | 12.25 | 1.23 | 1.58 | 3.17 | 0.87 | |||
I-18 | 18:04:49.7 | -30:03:02.5 | -0.039 | 16.73 | 15.06 | 13.79 | 13.03 | 12.96 | ||||
0.421 | 15.42 | 14.31 | 13.50 | 12.80 | 12.81 | 1.11 | 1.43 | 2.61 | 0.69 | |||
I-24 | 18:04:50.3 | -30:03:08.9 | -0.039 | 17.19 | 15.40 | 14.12 | 13.39 | 13.22 | ||||
0.421 | 15.88 | 14.65 | 13.83 | 13.16 | 13.07 | 1.23 | 1.58 | 2.81 | 0.76 | |||
I-36 | 18:04:50.9 | -30:03:48.0 | -0.026 | 16.39 | 14.42 | 12.90 | - | 11.86 | ||||
0.434 | 15.04 | 13.65 | 12.59 | 99.75 | 11.71 | 1.39 | 1.79 | 3.34 | 0.89 | |||
I-42 |
18:04:47.7 | -30:03:46.5 | -0.046 | 16.43 | 14.33 | 12.81 | 11.90 | 11.79 | ||||
0.414 | 15.15 | 13.60 | 12.53 | 11.68 | 11.64 | 1.55 | 1.99 | 3.50 | 0.89 |
Target | V | Date/ | UT | exp | Seeing | Airmass | ( S/N)/px | Slit |
![]() |
![]() |
Julian day | (s) | ('') | width |
![]() |
![]() |
|||||
I-16 | 17.496 | 26.06.00 | 03:46:28 | 5400 | 1.0 | 1.0 | 1.0'' | 208 | 206 | |
2 451 721 | 05:18:53 | 5400 | 0.8 | 1.0-1.2 | '' | '' | '' | |||
I-18 | 16.732 | 26.06.00 | 03:46:28 | 5400 | 1.0 | 1.0 | 40 | 1.0'' | 211 | 209 |
2 451 721 | 05:18:53 | 5400 | 0.8 | 1.0-1.2 | '' | '' | '' | |||
I-24 | 17.194 | 26.06.00 | 06:59:18 | 4584 | 1.1 | 1.2-1.6 | 30 | 1.0'' | 215 | 213 |
2 451 721 | 08:37:43 | 3600 | 1.3 | 1.7-2.6 | '' | '' | '' | |||
I-36 | 16.392 | 26.06.00 | 00:13:03 | 5400 | 1.1 | 1.2-1.8 | 40 | 1.0'' | 218 | 217 |
2 451 721 | 01:48:35 | 5400 | 1.1 | 1.2-1.0 | '' | '' | '' | |||
I-42 | 16.426 | 06.07.01 | 05:29:42 | 3600 | 0.7 | 1.0-1.1 | 30 | 0.8'' | 222 | 215 |
2 452 096 | 06:30:35 | 3600 | 0.8 | 1.2-1.5 | '' | '' | '' |
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Figure 1: Finding chart for the 5 stars observed in NGC 6528 (HST WFPC2 F814W image). Resolution is 0.1''/px. |
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