A&A 465, 185-196 (2007)
DOI: 10.1051/0004-6361:20066643
P. Sestito1,2 - S. Randich1 - A. Bragaglia3
1 - INAF-Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5,
50125 Firenze, Italy
2 -
INAF-Osservatorio Astronomico "G.S. Vaiana'' di Palermo, Piazza del Parlamento 1, 90134 Palermo, Italy
3 -
INAF-Osservatorio Astronomico di Bologna, Via C. Ranzani 1,
40127 Bologna, Italy
Received 26 October 2006 / Accepted 2 January 2007
Abstract
Context. We have carried out a big FLAMES survey of 10 Galactic open clusters aiming at different goals. One of them is the determination of chemical abundances, to put constraints on the radial metallicity gradient in the disk and its evolution. One of the sample clusters is the very metal-rich NGC 6253.
Aims. We have obtained UVES high resolution spectra of seven candidate cluster members (from the turn off up to the red clump) with the goal of determining the chemical composition of NGC 6253 and of investigating its origin and role in the interpretation of the radial metallicity gradient in the disk.
Methods. Equivalent width analysis and spectral synthesis were performed using MOOG and Kurucz model atmospheres.
Results. We derived abundances of Fe, -, and Fe-peak elements, the light element Na, and the s-process element Ba. Excluding two likely non-members and the clump giant, whose metallicity from equivalent widths is overestimated, we find an average [Fe/H] = +0.36
0.07 (rms) for the cluster. For most of the other elements we derive solar abundance ratios.
Key words: stars: abundances - Galaxy: evolution - Galaxy: disk - open clusters and associations: individual: NGC 6253 - Galaxy: abundances - open clusters and associations: general
The advent of new observational capabilities in the last years allowed
a steady improvement
in the field of chemical abundances in astrophysical objects
studied through high resolution spectroscopy.
In particular, it is now possible to derive precise element abundances
in old and distant Galactic open clusters, allowing us
to address different issues such as
the origin of the clusters themselves, the Galactic
radial metallicity gradient,
and the formation and evolution of the Galactic disk
(see, e.g., Friel 2006 and references therein).
At the
same time, the abundances of other
species, such as - and Fe-peak elements, or
the s- and r-process elements, and their ratios to
Fe, are crucial to obtain insights into the role of stars with different masses and
evolutionary lifetimes in the heavy element enrichment of the interstellar
medium. This,
through comparison with Galactic
enrichment models, permits us to put constraints
on the initial mass function and star formation history
during the early phases of disk evolution.
In this context, we carried out a VLT/FLAMES program
on a sample of 10 open clusters (Randich et al. 2005).
One of the main goals of this project is the determination
of the cluster metallicity and chemical composition
through the analysis of UVES spectra of evolved members (Sestito
et al. 2006 - hereafter
Paper I).
We focus here on the 3 Gyr old cluster
NGC 6253; this object, located
towards the Galactic center, is one of the most interesting
clusters
in our sample, since it has a metallicity considerably higher than solar.
Twarog et al. (2003)
suggested that
NGC 6253 might be the most metal-rich object in the Galaxy.
From photometric indices,
they found
[Fe/H
,
while from comparison with isochrones they concluded
that
-enhanced
isochrones provided the best fit, indicating a [Fe/H] closer to +0.4.
Besides the study by Twarog et al. (2003)
various photometric surveys of this cluster were carried out
in the last
10 years
(Bragaglia et al. 1997;
Piatti et al. 1998; Sagar et al. 2001);
most of them yielded
very similar values of reddening, age, and distance,
namely
,
age
3 Gyr,
and
(m-M)0=10.9-11.0(
kpc).
The first
spectroscopic determination of the metallicity of NGC 6253
was carried out by Carretta
et al. (2000), who found [Fe/H
.
A more recent analysis
(Carretta et al., submitted) based on better quality spectra
of five red clump stars favors a higher metallicity
([Fe/H]=+0.46).
Besides NGC 6253, a few other open clusters with metallicity higher than
solar are confirmed
by spectroscopic means (see Randich 2007 and references therein).
Among them we mention
NGC 6475 and Hyades
(ages 250 and 600 Myr, [Fe/H]=+0.14 and +0.13, respectively:
see Sestito et al. 2003 and Boesgaard & Friel 1990),
Praesepe (
600 Myr, [Fe/H]=+0.25; Pace et al., in preparation), the
1 Gyr old NGC 6134
([Fe/H]=+0.15, Carretta et al. 2004),
and the 2 Gyr old IC 4651 ([Fe/H]=+0.10: Pasquini
et al. 2004; Carretta et al. 2004).
The most noticeable
metal-rich cluster
is the very old (
8-9 Gyr) NGC 6791, recently investigated by
Origlia et al. (2006),
Carraro et al. (2006), and Gratton et al. (2006),
who derive metallicities of [Fe/H]=+0.35, +0.39, and +0.47, respectively
(from giant stars).
Table 1: Observation log of NGC 6253.
The origin of metal-rich disk clusters is puzzling, especially in the case of old ones, since the classical view of Galactic evolution predicts an over-time enrichment of the interstellar medium, and as a consequence only the youngest stars should have metallicities higher than solar. Nevertheless, other old and metal-rich stellar populations exist, such as bulge field stars (Fulbright et al. 2006, 2007) and objects in the solar neighborhood with kinematics and metallicities more similar to those of the bulge, in spite of their position (Castro et al. 1997; Pompeia et al. 2003). The high metallicity of stars in the center of the Galaxy can be explained with an early enrichment of that region, while the presence of old metal-rich stars/open clusters in the disk is more puzzling. One hypothesis about the origin of metal-rich open clusters is that they were born in the inner side of the Galaxy, close to the bulge, where the metal enrichment occurred early and rapidly, and then they moved outwards in the disk. Alternatively, they might have been born in an external environment and then captured by our Galaxy, as discussed by Carraro et al. (2006) for NGC 6791, although a very recent paper on this cluster excludes this possibility (Bedin et al. 2006). Finally, the simplest explanation is that metal-rich clusters originated in the disk itself, in a region characterized by faster enrichment. The element abundance distribution of metal-rich open clusters is very useful to put constraints on their origin.
Independently of their origin, the very existence of metal-rich open clusters provides ideal samples to investigate other topics, such as the dependence of light element depletion on chemical composition, or planet formation and evolution. In particular, it is now well ascertained that stars hosting giant planets are more metal-rich than stars not harboring planetary systems (e.g., Santos 2006 and references therein). Metal-rich clusters thus represent very good targets in which to search for planetary systems, although photometric searches for transiting planets have not been successful so far (but have shown feasibility - see, e.g., Paulson et al. 2004a,b for the Hyades; Mochejska et al. 2005 for NGC 6791).
We present here a new high resolution spectroscopic investigation of NGC 6253; with respect to Carretta et al. (2000, 2007, submitted) our sample covers a much wider region in the color-magnitude diagram (CMD), including not only a red clump member, but also turn off (TO) and subgiant/red giant branch (SGB, RGB) stars. Since, as we will show in the paper, at very high metallicities the Fe content of clump stars derived with equivalent width analysis might be overestimated - due to heavy line blending - the analysis of hotter TO and RGB stars should in principle provide more reliable results. The paper is organized as follows: in Sect. 2 we describe the sample and data reduction, while Sect. 3 is dedicated to the method of analysis and estimate of uncertainties. In Sect. 4 we report our results, checking the validity of the metallicity scale. The results are then discussed in Sect. 5, and summarized in Sect. 6.
The spectra of the NGC 6253 sample presented in this
paper were collected with FLAMES on VLT/UT2
(Pasquini et al. 2000), using the fiber link to UVES with
a spectral resolution of R=47 000.
The GIRAFFE fibers were instead used for collecting spectra of
a large number of main sequence stars, with the goal of
investigating lithium abundances and radial velocities (Randich et al.,
in preparation).
The observations were carried out
in service mode
during April 2004, with two FLAMES configurations,
using two different
gratings (CD3 and CD4, covering the wavelength ranges
4750-6800 Å and 6600-10 600 Å, respectively) for each of them.
We used two configurations to maximize the
number of objects observed with GIRAFFE; as a consequence,
the UVES pointings also changed. More in detail,
the two configurations differ for the
number of stars observed (seven and six stars were observed with UVES in configurations
A and B, respectively) and for the correspondence between fiber and object.
Since all the six stars observed with UVES
in configuration B are in common with configuration A, we observed seven stars in total.
Table 1 gives the log of observations for the cluster.
Data reduction was carried out by ESO personnel using the
dedicated pipeline, and we analyzed the 1-d, wavelength calibrated spectra
using standard IRAF
packages.
The contamination by atmospheric telluric lines was taken into account
by performing a correction on the spectra with the task TELLURIC in IRAF
via a comparison with early-type stars observed with UVES during another
run. Background subtraction was carried out, as customary,
using one fiber dedicated to the sky.
We report the target stars in Table 2, adopting the ID numbers from the EIS survey (Momany et al. 2001; Col. 1); since we used the photometry by Bragaglia et al. (1997; Cols. 6 and 7), we also list for completeness their IDs in Col. 2. The only star not included in the study by Bragaglia et al. (1997) is 105495, for which we adopted the EIS photometry calibrated to Bragaglia et al. (1997). As mentioned in Sect. 1, NGC 6253 was also investigated by Twarog et al. (2003); therefore we provide a cross-identification with their numbering system (Col. 3). The signal-to-noise (S/N) ratios reported in Col. 10 have been measured in the spectral regions around 5600 Å and 6300 Å.
Table 2:
Data for NGC 6253. ID
and
BV photometry (non corrected for reddening)
are from Bragaglia et al. (1997) except for
star 105495, see text;
we also report ID
,
used through out the paper,
and a cross-identification with Twarog et al. (2003).
The number of exposures
for each star is intended as a number of pointings
with the same cross-disperser.
We measured radial velocities (RV) with RVIDLINES using several tens of metallic
lines on each single spectrum, and subsequently we combined
multiple spectra. The heliocentric RVs (Col. 9 of Table 2) have
uncertainties of 1 km s-1with the exception of star 023501 for which the error is almost
3 km s-1.
The mean RV of the whole sample is
km s-1: RVs
of stars
069360 and 022182 deviate from this value
by more than 1
;
therefore we consider them non-members,
although we cannot exclude that they are binary cluster members.
If we exclude these two stars, we obtain
km s-1for the remaining 5 objects.
However, star 023501 also has an RV slightly deviating from the
two averages above (by
3 km s-1), therefore we will provisionally consider it
as a doubtful member (see below and Sect. 4).
By computing the average radial velocity considering only the
4 stars that can be safely classified as members,
we have
km s-1.
Figure 1 shows the CMD of
the cluster, where the selected stars are marked
using different symbols.
Note that
the positions in the CMD of
stars with radial velocity deviating from the mean
are consistent with membership (in particular
that of the doubtful member 023501).
Notes on the evolutionary status and membership of the stars are shown in
Col. 11 of Table 2.
Finally, in Fig. 2 we show the spectra
(referred to RV=0) in the wavelength region around the H
feature
for all the observed stars.
![]() |
Figure 1: Color-magnitude diagram for NGC 6253. The observed stars are evidenced by circles (members), triangles (the radial velocity doubtful member, but with metallicity consistent with membership), and squares (non-members). |
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The analysis of chemical abundances was performed
by means of equivalent widths (EWs) using
an updated version (2006) of the package MOOG (Sneden
1973) and using model
atmospheres by Kurucz (1993). MOOG works under the assumption
of local thermodynamic equilibrium (LTE).
Solar abundances of Fe and other elements (Na, Mg, Si, Ca, Ti,
Cr, Ni, Ba) were derived to determine the
zero point of the metallicity scale.
The line lists adopted for the Sun and for evolved stars
were retrieved from
Gratton et al. (2003) and
are described in
Paper I and
Bragaglia et al. (in preparation).
We recall from Paper I
that when available we adopted
collisional damping coefficients from Barklem
et al. (2000),
otherwise we considered
the coefficients by Gratton et al. (2003), or,
for a few lines, the classical Unsöld (1955) approximation.
For the Sun we obtain
I
(standard deviation, or rms) using
K,
,
and
km s-1.
![]() |
Figure 2: NGC 6253 sample spectra in the wavelength region at 6500-6600 Å. |
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Table 4: Stellar parameters and Fe abundances for stars in NGC 6253. Numbers in parenthesis in Col. 6 are [Fe/H] estimated by spectral synthesis.
The spectra were normalized using CONTINUUM in IRAF, dividing the spectra in small regions (50 Å) and visually checking the output; EWs for the various lines were measured with SPECTRE (developed and maintained by C. Sneden) by Gaussian fitting of the line profiles. We provide the EWs for each star in Table 3 (available only in electronic format the CDS): the first two columns list the wavelength and the element, and the others show the corresponding EW for each star. Note that the definition of the continuum and the determination of the EWs are a very critical step for our stars, since they are rather cool objects with low gravity and with an exceptionally high metallicity. As a consequence, several lines can be affected by strong blending (see below also).
Initial effective temperatures (
)
and gravities were
estimated from
photometry. In the case of giant stars we used the
B-V vs.
calibration by Alonso, Arribas, & Martinez-Roger
(1999) for giants, while
for TO stars we adopted the calibration by
Alonso et al. (1996),
based on a large sample of dwarfs.
Surface gravities were
derived using the
expression
,
where M is the mass and
the bolometric magnitude (with
). Bolometric corrections for giants
were derived following Alonso et al. (1999), while
those for the two TO stars were retrieved from Johnson (1966)
and are close to 0.
We adopted the most recent cluster parameters by
Bragaglia & Tosi (2006),
(m-M)0=11.0,
E(B-V)=0.23, and an age of 3 Gyr, which,
using Z=0.05 isochrones, corresponds to masses
at the TO and
1.40
at the clump. We assumed
for all the stars, since uncertainties
0.1
translate into differences of 0.03 dex or lower in surface gravities, which
are well below the random errors and do not affect the metallicities
derived from Fe I lines.
Note that the calibrations by Alonso et al. (1996; 1999)
are valid up to a metallicity
of [Fe/H]=+0.2, while the cluster should have a higher Fe content;
nevertheless,
the weak dependence of the
on B-V colors suggests that the error
is small, and in any case the photometric temperatures are
used only
as initial parameters and are optimized during the spectroscopic analysis.
In Paper I
the microturbulent velocities were derived
using the relationship by Carretta et al. (2004),
based on the
optimization of Fe abundance as a function of theoretically expected
EWs for the given
lines (see the quoted references for further details).
The formula cannot be safely applied for NGC 6253, since it was
based
on a sample of giants in clusters with
metallicities closer to solar; in this case we deal instead with a very metal-rich cluster and also
with TO stars. Therefore, we only used the microturbulences
by Carretta et al. (2004) as starting values; then, we optimized them
by minimizing the slope of
(Fe) I vs. the observed EWs; when
plotting Fe abundances as a function of expected EWs, we obtain
a rather small trend (i.e., the slope of
the relationship is smaller than its error), suggesting that the two methods of analysis
are in fair agreement.
As is customary, final effective temperatures were derived during the analysis
(after 1
clipping) by minimizing the trend
of
(Fe I) vs. the excitation potential (EP);
as for surface gravities,
we cannot derive them from
the ionization equilibrium condition (i.e., the assumption that
the difference between Fe I and Fe II abundances
in the stars analyzed
should be similar to that found for the Sun)
since the few lines of
ionized Fe are strongly affected by blending.
On the other hand, since in most cases
spectroscopic temperatures are in good
agreement with photometric ones and the distance
is well known, we decided to adopt the
photometric
values and left them unchanged during the analysis.
The only stars for which the spectroscopic and photometric
differ by
a rather significant amount
(
200 K) are the TO stars;
note, however, that changing the
by
200 K
would imply a variation of
dex in
,
consistent with
the errors (see Sect. 3.3).
For the reasons mentioned above, we retained only
neutral Fe features in the analysis.
In Table 4 we list
the photometric
and spectroscopic
(Cols. 2 and 4),
the adopted surface gravity (photometric, Col. 3), and
the spectroscopic microturbulence (Col. 5) for each star.
The major source of random
uncertainties affecting element abundances
derive from errors in EWs and uncertainties
in stellar parameters; systematic uncertainties
come from biases due to the method of analysis
adopted
and from errors in the line list (i.e., possible blendings and
oscillator strengths).
The errors in abundances related to EWs are given in good approximation
by the standard deviation (rms) around the mean abundance derived
from individual lines for each star, call it .
This rms also includes the errors related to atomic parameters:
values were taken from the literature, thus
we cannot give a precise estimate of the effect of their uncertainties;
however, since
our abundance scale is directly related to solar abundances
and the line list used for cluster stars is
very similar to that adopted for the Sun, we can assume that
internal errors due to uncertainties in the atomic parameters
are minimized.
Note that when the abundances of elements other than Fe
are expressed as [X/Fe]
, a total
should be computed by quadratically adding the rms
on [X/H] and on [Fe/H].
Table 5:
Sensitivities of abundances
((X)) to variations in the atmospheric parameters for
TO and clump stars in NGC 6253.
The contribution of uncertainties in stellar parameters,
,
,
and
,
were estimated by varying each parameter of a given quantity (leaving
the other two unchanged) and then adding the three errors.
We assumed variations of
70 K in
and
0.10 km s-1in
,
since these changes would introduce significant trends into
the relationship of Fe abundances with EP and observed EWs.
We cannot estimate an uncertainty in
in a similar way ,
since we did not
optimize the gravity using the ionization equilibrium.
However, errors in
are usually of the order
of
0.15-0.25 dex (e.g., Paper I);
thus, we assumed a conservative
dex.
Table 5 shows the sensitivity of elemental abundance
(
(X))
to variations in the atmospheric parameters for two cluster members:
the TO star 069885 and the clump star 105495.
In the case of Fe we computed
,
the quadratic sum of the three errors due to uncertainties in the
stellar parameters; this was not done for the other elements, since
in the final computation of [X/Fe]
one should take into account the
for [X/H] and for [Fe/H], which
could go in opposite directions (see Table 5).
Finally, we wish to give an estimate of the systematic uncertainties
in the Fe abundance scale related to
the method of analysis. This can be done for example by
analyzing a star with a well-known metallicity, possibly observed
with the same instrument and
using the same method
of analysis.
Since we did not collect spectra of stars outside of the clusters
included in the program,
we performed the analysis for two clump stars in the Hyades
observed with SARG at TNG at similar resolution.
A detailed description of the analysis of the two Hyades
is reported in Paper I; here, we only mention that we did find
a [Fe/H] in reasonable
agreement with the literature estimates, confirming that,
up to the metallicity of the Hyades,
our method of analysis should not be affected by large systematic errors.
To check if the metallicity scale is also correct for
a very high metal content, we carried out some tests on
Leonis (see Sect. 4.3), a rather
luminous giant that is known to have a
remarkably oversolar metallicity (e.g., Gratton & Sneden 1990).
We anticipate here that at the metallicity of NGC 6253 our Fe
abundance for the coolest star
(clump) might be overestimated by
0.1-0.15 dex.
Fe abundances are listed in
Table 4, together with their errors, in Cols. 6-8:
,
the standard deviation
from the mean abundance obtained over the whole set of lines for each star,
and
,
the
total uncertainty in which we also consider errors
due to stellar parameters
(
).
Star 023501 was classified as a doubtful member
from its radial velocity (see Sect. 2): we found
[Fe/H]=+0.29 for it, in agreement with those of the confirmed members, thus we
conclude that this object is a probable cluster member.
Stars 069360
and 022182 are instead radial velocity non-members but, as
already mentioned, we cannot exclude that they are binary cluster
members.
The first one
has a high Fe abundance (+0.48)
similar to those of members; nevertheless,
since we do not optimize gravities from the ionization equilibrium
and we rely on the photometric
values, the derived [Fe/H] value might be due only to
a coincidence; in other words, if the adopted distance is
wrong, one finds a wrong metallicity. On the other hand, if the
star
would effectively be a binary cluster member, the [Fe/H] we found
might be the correct one.
Star 022182
has a much lower [Fe/H] with respect to other stars,
that is +0.12.
In any case, we disregard the two non-members in the following.
The average metallicity (computed excluding
the non-members) with the rms error
is also shown in Table 4, [Fe/H
.
For the 4 ascertained members we find [Fe/H]=+0.45 (TO star),
+0.49 (clump), and +0.32, +0.39 (SGB/RGB).
The metallicity of the clump star 105495 based on EWs could likely
be overestimated, due to unresolved blends (the spectrum is
very crowded due to the combination
of high Fe content and low temperature);
by excluding the clump star, the average
metallicity
slightly decreases to [Fe/H
.
A possible offset in the metallicity scale will
be discussed in the next sections (4.2 and 4.3).
Values shown in parenthesis in Col. 6
of Table 4 are the metallicities found from spectral synthesis
for the hottest TO star and for the clump star;
the average (+0.34) computed taking into account these values is also reported.
Figure 3 shows [Fe/H] values as a function of
effective temperature for all the stars observed.
The average cluster metallicity
rms ([Fe/H
)
are indicated.
![]() |
Figure 3:
Fe abundances as a function of
![]() |
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To check Fe abundances derived through the EW analysis, we
carried out spectral synthesis for the warmest (069885) and
coolest (105495) cluster stars,
which are also those having the highest [Fe/H].
The spectral synthesis was performed in
a spectral interval of
Å around the Li I 6707.8 Å line.
As for the EW analysis, we used MOOG and Kurucz atmospheres; an earlier
version of MOOG (2000) was, however, employed, with
a line list optimized to fit the solar spectrum obtained
with UVES; the Fe abundances retrieved from the synthesis are therefore
differential with respect to the Sun.
We used
the classical
Unsöld (1955) approximation for the damping coefficients,
since the spectral range investigated with the synthesis does not
include strong lines.
Synthetic spectra were computed adopting stellar parameters derived
by the EW analysis and listed in Table 4.
For each star, we computed
a synthetic spectrum with the metallicities determined through EWs and
then others until the best fit of the observed spectrum was obtained.
Figure 4 shows the spectral synthesis
for two stars (105495 and 069885) in the wavelength range
6700-6718 Å.
In the upper panel, we plot the comparison between the observed spectrum of
the clump star
and two synthetic spectra with [Fe/H]=+0.35 and +0.49:
as is evident, the lines of Fe I
in the synthetic spectrum with metallicity +0.49 are in the majority of cases deeper
than the observed ones, while the synthesis
with [Fe/H]=+0.35 provides a better fit (although far from perfect).
This also happens for the TO star (lower panel)
for which we computed two synthesis with [Fe/H]=+0.35 and +0.45.
Therefore, in both cases the metallicity obtained from
spectral synthesis is dex
below that determined using EWs.
A higher metallicity from EW analysis
with respect to synthesis
can easily be explained for
the clump star, whose spectrum
might be affected by blending; on the other hand,
the discrepancy found for the TO star is rather
surprising, and we do not really have an explanation for this.
![]() |
Figure 4: Comparison between observed (solid line) and synthetic spectra for the clump star 105495 ( upper panel) and the TO star 069885 ( lower panel) in the Li I region (6707.8 Å). The (red) thin dots are spectra computed adopting the spectroscopic metallicities (+0.49 for 105495 and +0.45 for 069885), while the (blue) thick dots are spectra with [Fe/H]=+0.35. |
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As mentioned in Sect. 4.1, by assuming [Fe/H]=+0.35for the TO and clump stars, the average [Fe/H] of the cluster would decrease down to +0.34. However, we stress that errors in the determination of the continuum in the observed spectra also affect the comparison with synthetic spectra and not only EW measurements. Finally, note that in the wavelength region shown in Fig. 4 there are observed spectral features not reproduced by the synthesis: these are lines of elements other than Fe, which in this case are not important for the determination of the metallicity (e.g., the Ti I line at 6706.3 Å and the Li I line at 6707.8 Å) and for which, therefore, we did not optimize atomic parameters and abundances.
Since the analysis through EWs and the spectral synthesis yield
slightly different results, we
carried out a further test
to check our metallicity scale: namely, we analyzed
the metal-rich giant star
Leonis.
The values of [Fe/H] estimated in the literature are all around +0.3-0.4 dex
(e.g., Gratton & Sneden 1990, [Fe/H]=+0.34; Fulbright et al.
2006, [Fe/H]=+0.32; Gratton et al. 2006,
[Fe/H]=+0.38), therefore similar to that of NGC 6253.
We analyzed a spectrum of
Leo
observed with FEROS on the 2.2 m telescope at La Silla Observatory
with a resolution similar to that of our sample stars.
By using the same line list as for NGC 6253, we obtained [Fe/H]=+0.51 for
Leo, and
K,
,
and
km s-1; whereas
the atmospheric parameters are in good agreement
with the determinations of other authors,
the Fe abundance is higher than previously found.
A metallicity more similar to those quoted by other authors is obtained by us
with the spectral synthesis, i.e., [Fe/H]=+0.38;
since by spectral synthesis we find [Fe/H]=+0.35 for 105495, which has atmospheric
parameters similar to those of
Leo, our results
suggest that the clump star in NGC 6253 and
Leo should actually
have similar metallicities, but a scale offset is present between the
synthesis and EW analysis.
Figure 5 shows a comparison between
the spectra of
Leo and 105495
in the wavelength region around the Li I doublet at 6708 Å,
where several Fe I features are present:
as is clearly visible, the spectral lines of the two stars are similar,
suggesting
that the metallicities are nearly the same. In particular,
Leo is slightly
colder than 105495, and indeed its metal lines
are slightly stronger.
In any case,
the problem of the determination of
a zero-point for the [Fe/H] scale remains.
To further check the metallicity
scale and the origin of the discrepancy,
we carried out different tests on
Leo.
We considered the line list for Fe I adopted by Fulbright
et al. (2006): more in detail, using the 30 lines in common
with our list we carried out an EW analysis adopting our atomic
parameters and their EWs.
The analysis by Fulbright et al. (2006)
is differential with respect to the Sun, for which
they derive
(Fe
.
With their measurements and our atomic parameters we find [Fe/H]=+0.40,
i.e., 0.08 dex larger than that of Fulbright et al. (2006;
+0.32).
Note that we obtain
K, higher
than the value we previously found, but in agreement
with Fulbright et al. (2006).
If we adopt the 30 lines in common with Fulbright et al. (2006), but using our EWs, we find a metallicity [Fe/H]=+0.49,
with
K.
Finally, we also repeated the latter analysis for the clump star
105495 in NGC 6253, i.e., with the lines in common with
Fulbright et al. (2006 - and obviously our EWs), and
we obtained [Fe/H]=+0.44.
Part of the discrepancies with the previous analysis can be due
to EW measurements (program/method and continuum tracing)
and the adopted code, but in any case we obtain a systematically higher
metallicity.
Given the disagreement found with the literature results
for
Leo, and between the spectral synthesis and EW analysis for
stars in NGC 6253,
we conclude that
an offset
in the abundance of the clump star is present
(probably due to a combination of low
and high [Fe/H]
that results in very strong and blended features); whereas
we are not able to precisely quantify the offset,
we note that it is in the range
+0.10-0.15 dex;
for the other stars we cannot estimate a possible abundance shift
through a similar comparison with a well-known star,
since we do not know a very metal-rich subgiant to be used
as a reference object.
In summary, we will adopt the [Fe/H] values derived
through EW analysis for all the stars,
but with the caveat that those of the clump
and of the hottest TO stars are probably overestimated,
while
the metallicity found for the remaining stars is likely to be correct.
![]() |
Figure 5:
Comparison between the spectra of the
clump star 105495 in NGC 6253
(solid line) and of ![]() |
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As far as other elements are concerned, it has not been possible to
carry out a direct comparison with
Leo.
Although a chemical analysis of this star has
recently been performed by Fulbright et al. (2007),
their line lists and ours have only a small number of
lines (if any) in common for each element; moreover, their
EWs were not published.
By comparing the abundances
for
-elements (Ca and Si in particular) obtained using
the few lines in common with Fulbright et al. (2007)
and our EWs, we obtain [X/H] values
similar to theirs, suggesting that the scale offset is likely to
affect only Fe.
Table 6:
Elemental ratios ([X/H]) for stars in NGC 6253: Si, Ca, Ti, Cr, Ni, and Ba.
Errors are the rms
- due to
EW uncertainties - on [X/H]. The solar abundances we found
are also shown (
(X)
).
Table 7:
[X/Fe] abundances and averages.
Errors are the quadratic sum of
on [X/H] and on [X/Fe].
Table 8: Abundances of Mg and Na (LTE and non-LTE) computed for each line adopted.
![]() |
Figure 6:
[X/Fe] abundances of the various elements analyzed
as a function of [Fe/H]. The solid and dotted lines
are the averages ![]() |
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Besides Fe, we derived the abundances of the light isotope
Na, the -elements Mg, Si, Ca, Ti,
the Fe-peak elements Cr and Ni, and the s-process element Ba.
[X/H] values for Si, Ca, Ti, Cr, Ni, and Ba
are listed in Table 6, together
with the errors due to uncertainties in EWs (
).
The
(X) values in the Sun found by us (with the exception of Ba), are also shown (last column).
[X/Fe] ratios are instead shown in Table 7,
with errors
computed by quadratically adding the rms
for [Fe/H] and for [X/H].
We notice that for Ba and Si
the measurement is based on
a small sample of lines (3 and 4-6 features,
respectively).
For the other species (Ca, Ti, Cr, and Ni), we carried out the analysis
using a
large set of lines, and we performed 1
clipping.
In the last three columns of Table
7 we report
two average values (
rms) of the [X/Fe] abundances for the cluster.
The means
are computed including all the five stars (Average1),
or
excluding the clump star (Average2).
In general the two averages are consistent
with each other within the errors; the scatter is smaller if only the confirmed members are considered.
As discussed in further detail by Bragaglia et al. (in preparation),
we find a rather large error (0.14)
on the solar
,
and this might
depend on uncertainties on the
s and on the fact that
the analysis of this species is based on a very small number of lines.
The analysis of Na usually is based on a set of seven lines, but
in the case of NGC 6253 not all these features are measurable
and in addition the few lines used give discrepant results.
For this reason, the ratios of Na and Mg to Fe in stars
of NGC 6253 have been computed
from a line-to-line comparison, instead of comparing the average
abundance of a star to the solar one.
The abundances of the two elements are presented separately in Table 8,
where we report [X/H] and [X/Fe] for each line
in common between the lists adopted for the Sun and for the cluster stars.
We mention that
for Mg only two lines were used, and, as a consequence,
the results shown in Table 8,
which indicate a Mg enhancement,
should be taken with caution and would require a dedicated study
that goes beyond the primary goal of the present paper.
We used three spectral features
for the determination of Na abundance, but the two lines at 6154-6160 Å are likely
to be more reliable than the 5688 Å line, which is rather strong and might
deserve more detailed damping computations.
Na abundances derived with MOOG are based on the assumption of LTE.
As is well known, this assumption may introduce systematic errors
in the computation of the abundance;
whereas for most of the elements it has been ascertained that
non-LTE corrections are negligible,
in the case of Na they might be important.
The problem is that non-LTE effects affect
stars in diverse evolutionary phases (main sequence, TO, RGB, clump)
at different levels
since they are strongly dependent on the temperature and surface gravity.
Also, discrepant results have been obtained from different authors
in the computation of non-LTE corrections.
For example, Gratton et al. (1999)
find moderate negative corrections of the order of
0.05-0.1 dex for giant stars, while Mashonkina
et al. (2000) estimate larger corrections, of the order of
0.15
dex.
In Table 8 we show [Na/Fe] values derived with MOOG
and corrected adopting the tabulations by Mashonkina et al. (2000).
Considering LTE abundances, Na seems to be enhanced
with respect to the solar value, in agreement with other findings for open clusters
(Friel
et al. 2003; Yong et al. 2005;
Bragaglia et al. 2006), but a certain amount
of scatter is present in
the abundances from the various lines and also among
the various stars.
On the other hand, with the non-LTE corrections by Mashonkina et
al. (2000) the [Na/Fe]
values end up being lower by
0.10-0.15 dex depending on
the line considered and
on the stellar parameters. In this case, the average Na abundance of the
cluster would be nearly or slightly above solar;
therefore, we suggest that the Na abundance enhancement
claimed for open clusters based on giant stars might
be in part related to non-LTE effects (see also Randich et al. 2006).
On the other hand, the [Na/Fe] abundance ratios of field dwarfs
do not appear to be enhanced (e.g., Soubiran & Girard 2005);
this issue
deserves further investigation, which is beyond the goals of
this paper.
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Figure 7:
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Figure 8:
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![]() |
Figure 9:
![]() |
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In Fig. 6 we show the element abundances of Table 7
as a function of [Fe/H];
also in this case the mean abundance
rms are represented.
We adopted the value Average1 shown Table 7,
i.e., computed including all the 5 stars.
All the elements, apart from Ba, have average solar abundances;
Ba is enhanced,
as already found for other clusters, e.g., by our group (Bragaglia et al.,
in preparation)
or by Bragaglia et al. (2006); however, abundance for this
element has been found to vary significantly between clusters (e.g., Gratton et al. 2004).
As is mentioned in the introduction,
the metallicity of NGC 6253
is unusual for disk stars,
which normally have abundances close to solar.
The only other cluster for which a very high Fe content has been reported
is the 9 Gyr old NGC 6791.
Carraro et al. (2006) and Origlia
et al. (2006) found [Fe/H]=+0.39 and +0.35, respectively,
and solar -element abundances.
NGC 6791 is a peculiar cluster, since it is
very massive and has a very eccentric
orbit; therefore, it is different
from NGC 6253 in several aspects.
In particular,
NGC 6253 is much younger than NGC 6791
and much less massive;
the orbit of NGC 6253 has not been studied yet, but it could be
interesting to have information on it.
Given its position towards the
Galactic center, NGC 6253
could have been born either towards the bulge, where
the metallicity is high, or in a region of the disk
where a particular metal enrichment occurred.
To get insights on the origin of this cluster,
in Fig. 7 we show
a comparison for
-element
abundances (Si, Ca, Ti) vs. [Fe/H]
in NGC 6253, NGC 6791,
and other open clusters with [Fe/H
;
the clusters are
the Hyades,
NGC 5822,
IC 4651, IC 4725, and
NGC 6705
(references can be found in Friel 2006 and
Randich 2007).
We plotted these three elements since their analysis is based
on a rather large sample of lines with respect to the other species,
and they are all explosive nucleosynthesis
-elements, i.e.,
they originate from type II Supernova events.
Note that for NGC 6253 we consider here
(and in Figs. 8 and 9) the values labeled as
Average2 in Table 7, excluding the
clump star; similarly, for Fe we adopted the value +0.36, again
excluding the clump star.
The figure shows that the average
-element abundances of
NGC 6791 and NGC 6253 are identical or in very good agreement.
-element abundances in open clusters with
oversolar Fe content are in general
close to solar
(with the exception
of a Si enhancement in IC 4725, Luck et al. 1994).
Figure 8 shows a comparison between NGC 6253 and disk dwarfs
observed by Mishenina et al. (2004; thin and thick disk)
and Bensby et al. (2005; thick disk). Also in this case
the -elements Si, Ca and Ti were considered;
the
figure shows the range in metallicity
[Fe/H
.
The average abundances of evolved stars
in NGC 6253 match the general
trend observed for thin and thick disk dwarfs very well.
Finally, in Fig. 9
we plot a comparison between NGC 6253 and
the results for bulge giant stars recently analyzed by
Fulbright et al. (2007)
and bulge-like field stars by
Castro et al. (1997) and
Pompeia et al. (2003). In the latter works
samples of nearby dwarfs
with kinematics and metallicity characteristics of a probable inner
disk or bulge origin have been investigated.
The bulge and bulge-like
field star samples cover a [Fe/H] range
,
but
we show abundances only for [Fe/H] larger than
-0.1.
Bulge and bulge-like stars are characterized by
-enhancement
at very low metallicities (not visible in the figure)
with
a decrease towards solar and oversolar metallicities, as shown
in the plot;
note, however, the much
larger dispersion with respect to disk stars reported in Fig. 8:
indeed there are stars showing
enhanced
-element abundances at solar/oversolar metallicities also.
![]() |
Figure 10: Radial gradient ([Fe/H] vs. Galactocentric distance) for open clusters. The results for clusters in our sample analyzed so far (filled circles; this paper, Paper I and Bragaglia et al., in preparation) are compared to other clusters analyzed with high-resolution spectroscopy (open circles, see Paper I for references) and low-resolution spectroscopy (asterisks, Friel et al. 2002). |
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The comparisons between NGC 6253 and field stars suggest that
the abundance of the cluster is in good agreement with the trend
observed in the disk (Fig. 8).
From this evidence,
we can speculate that NGC 6253 was born in the Galactic disk,
in a region where a larger than normal Fe enrichment occurred; on
the other hand,
we do not observe an enhancement of -elements with
respect to Fe.
Under the assumption that the cluster formed in the disk,
we can use it for the determination of the radial metallicity
gradient.
In Fig. 10 we show [Fe/H] as a function of the
Galactocentric radius for NGC 6253 and the other samples included in
our program (Paper I and Bragaglia et al., in
preparation),
compared to other samples analyzed
with high resolution spectroscopy (references for all the clusters investigated at high resolution can be found
in Paper I).
In the figure, we also show the low-resolution sample by
Friel et al. (2002; but note that we
excluded
the clusters in common with high resolution studies).
For consistency, we adopt for all the clusters
from Friel (1995) and Friel et al. (2002).
NGC 6791 (the cluster with the highest [Fe/H])
lies above the average trend for open clusters,
confirming that it
might have had an origin and evolutionary history
very different from those of other clusters (Carraro et al. 2006).
Also the very metal-rich NGC 6253 (the other
cluster with [Fe/H]>+0.3) and
Praesepe ([Fe/H]=+0.25;
Pace et al., in preparation) might lie above the mean trend.
In other words, the inclusion of
NGC 6253 in the [Fe/H] vs.
distribution makes the negative slope
of the gradient steeper.
However, we note that all the clusters with
lower than that
of the Sun -
8.5 kpc - have higher
than solar metallicities, with the exception of Cr 261,
which has [Fe/H]=-0.03(Carretta et al. 2005; De Silva et al. 2006). This is true for
high resolution spectroscopy results, whereas Friel et al. (2002),
from low resolution data, quote
metallicities from solar down to -0.25 dex for clusters
with
between 8 and 8.5 kpc.
Indeed, the metallicities by Friel et al. (2002) from low
resolution data are - for all the clusters and at all
s
- lower than those obtained from high resolution
analysis.
Considering only [Fe/H] values from high resolution and clusters
with
kpc,
we note that Cr 261 might represent an exception for its relatively low
Fe content, rather than NGC 6253 and other high metallicity clusters.
To our knowledge, no investigations of the orbit of
Cr 261 are present in the literature, but it would be very interesting
to have such information, to understand if this cluster might have
formed at a larger Galactocentric
distance than its present position.
We tentatively conclude that the problems/questions raised by
the very high metallicity of NGC 6791
should not really concern NGC 6253,
which is rather young and
is located towards the Galactic center and very close to it
(at variance with NGC 6791, which is
is much older and is not located in the Galactic center direction).
Finally, we would like to remark that, as mentioned in Sect. 1,
Twarog et al. (2003) suggested the possibility of an -element
enhancement in NGC 6253. We find instead solar-scaled abundances
for these elements,
implying that
stellar evolutionary models at these metallicities also need to be improved.
We report on chemical abundances in the metal-rich cluster NGC 6253, observed with VLT/FLAMES. The original sample includes seven stars (two at the turn off, one at the clump, and four on the subgiant/red giant branch). We find the following results:
Acknowledgements
P.S. acknowledges support by the Italian MIUR, under PRIN 20040228979-001, and by INAF-Osservatorio Astrofisico di Arcetri, where this work was completed. We thank the anonymous referee for her/his valuable suggestions. We are grateful to C. Sneden and S. Lucatello for having provided an updated version of MOOG and for useful discussion about it; we also thank G. Carraro and L. Pasquini for helpful discussion and suggestions.