A&A 378, 455-465 (2001)
DOI: 10.1051/0004-6361:20011210
V. G. Klochkova1,2 - N. N. Samus3,4
1 - Special Astrophysical Observatory,
N. Arkhyz 369167, Russia
2 -
Isaac Newton Institute of Chile, SAO Branch, Russia
3 -
Institute of Astronomy of Russian Acad. Sci., 48,
Pyatnitskaya Str., Moscow 109017, Russia and
Sternberg Astronomical
Institute of Moscow University, 13, University Ave., Moscow 119899,
Russia
4 -
Isaac Newton Institute of Chile, Moscow Branch, Russia
Received 27 June 2001 / Accepted 20 August 2001
Abstract
CCD spectra obtained with the echelle spectrometer of
the 6-meter telescope were used to determine, by the model
atmospheres method, the fundamental parameters
(=4800K, logg = 0.7) and detailed chemical
abundances for the star K413, a member of the globular cluster
M12. The resulting value,
,
is
the first metallicity determination for M12 using
high-resolution spectra. The main characteristic feature
of the star's atmospheric chemical abundance pattern is a large oxygen
excess,
.
The s-process elements are probably slightly depleted compared to
metallicity:
for yttrium and
zirconium,
for barium. Abundances
of the heavier elements: La, Ce, Nd, and Pr, do not differ,
relative to iron, from the solar ones:
.
The europium excess,
,
is
typical of members of low-metallicity globular clusters.
The spectrum of K413 shows, for the H
line, an variable
absorption and emission profile. From its high luminosity and
chemical abundance anomalies, we can suppose that K413 is in an
evolutionary stage after the AGB. In the spectrum, we
find absorption details that can be identified with diffuse
interstellar bands displaced by 16kms-1 to longer wavelengths
relative to the star's velocity.
Key words: globular clusters: general - globular clusters: individual M12 - stars: abundances - stars: atmospheres - stars: AGB and post-AGB
This paper continues the series of our publications on the studies of stars at a short-lived evolutionary stage of the transition from the asymptotic giant branch (AGB) to planetary nebulae (cf. Klochkova et al. 2000 and references therein). The main goal of this program is to investigate anomalies of the chemical composition for the objects that have experienced a sequence of changes of their energy sources, mixing, and matter dredge-up from layers with modified chemical composition to the surface.
Stars at late stages of evolution, members of globular clusters, are of particular interest: their cluster membership makes their evolutionary stage and luminosity known more confidently than for stars in the galactic field. Simultaneously, globular cluster membership makes high spectral resolution observations of a star difficult because such clusters are very distant from the Sun. Besides, it is difficult to prove membership, especially for stars above the horizontal branch. In fact, only one of the generally-used criteria, the value of the radial velocity, is decisive for such objects. For this reason, accurate radial velocity measurement is a special task.
In this paper, we present the results of our spectroscopic study
of a faint star (
), K413 (the number
according to Küstner 1933), in the globular cluster M12. In
Racine (1971), the star has the number II-02-51; Zinn
et al. (1972) list it as No.8. Note that Zinn et al.
(1972) compiled a large list of globular-cluster stars
above horizontal branches in the color-magnitude diagrams and
suggested to call them "UV-bright stars''. The Harris
(1996) catalog contains the following principal
characteristics of this cluster: the mean radial velocity,
kms-1 (relative to the local standard of rest,
kms-1); the mean color excess,
E(B-V) = 0.19; the distance modulus,
;
the distance, d = 4.9kpc.
The cluster's modern
V - (B-V) diagram can be found
in Brocato et al. (1996).
The diagram shows that K413 is brighter than AGB stars with
the same (B-V) color index by one magnitude, so it can
be considered as a post-AGB star.
The high probability of the cluster membership for K413 follows
from proper motions of stars in the field of the cluster M12
(Geffert et al. 1991) as well as from radial velocity data.
From moderate-resolution spectra, Harris et al. (1983)
measured radial velocities for a sample of such stars and were
able to confirm membership of several UV-bright stars in the
globular clusters M12 and M56 (in particular, the membership
of K413 in M12). The latter authors give the
mean velocity value
kms-1 for stars in the
cluster M12.
For K413, they derived the velocity
kms-1.
Our paper is the first presentation of results based upon high-resolution spectra of objects in the cluster M12. Section 2 briefly describes the techniques of observations, reductions, and analysis. Section 3 presents the results of our determination of the detailed atmospheric chemical composition of K413. Section 4 contains discussion and main conclusions.
Two echelle spectra in the 4300-7800Å range were obtained
in the prime focus of the 6 m telescope, with the echelle
spectrometer PFES (Panchuk et al. 1998).
The mean times of observations are collected in Table 1.
PFES is equipped with a
pixel CCD, designed
at the Special Astrophysical Observatory; the pixel size is
m. The spectrometer's resolution is up to 17000 in a wide wavelength range. The spectrum of a Th-Ar lamp
was recorded for wavelength calibration. Processing of the 2D
images (standard procedures of dark current subtraction, removal
of cosmic rays, subtraction of scattered light, extraction of
echelle orders) made use of the ECHELLE context of the MIDAS
system (version 1998). Spectrophotometric and positional
measurements of 1D spectra used the DECH20 package (Klochkova
& Galazutdinov 1991; Galazutdinov 1992). The
signal-to-noise ratio of the spectral portions used for the
chemical abundance analysis considerably exceeded 100, enabling
us to measure equivalent widths of weak lines, to 10-15 mÅ.
Note that results presented here are based mainly on the spectrum
s27605, while the second spectrum s31208 was obtained and used for
comparison only.
Spectrum | JD |
![]() |
||||||
Metals (n) |
![]() |
Na D1, D2 | DIB (n) | |||||
em, blue | abs | em, red | Stellar | Circumstellar | ||||
s27605 | 2451708.378 | -40.96 (453) | -88.2 | -43.8 | +17.4 | -38.6 (2) | -24.6 (2) | -24.9 (8) |
s31208 | 2452072.403 | -41.48 (125) | -106.9 | -53.4 | +16.0 | -40.91 (2) | -30.061 (2) | -23.58 (4) |
To derive the basic parameters of the stellar model atmospheres:
effective temperature,
,
and surface gravity,
logg, for computations of chemical abundances and of
synthetic spectra, we used the grid of model atmospheres
computed by Gustafsson et al. (1975) in the hydrostatic
approximation.
The wide spectral range simultaneously recorded by the PFES
echelle spectrometer makes it possible to determine the principal
model parameters,
and logg, using only spectroscopic
criteria free of the influence of interstellar and
circumstellar reddening.
The "photometric'' criteria are used widely for
temperature determination, especially for metal poor stars. But
this method is not useful for the star K413. For objects at an
advanced evolutionary stage it is specially difficult to account
for reddening. We do not know a priori normal (intrinsic) colours
for such an object of uncertain evolution any stage and uncertain
chemical composition. Using average "
-color''
calibrations and average colour-excess, we do not take into
account peculiarities of atmospheres of individual stars. This,
in turn, may lead to essential errors in temperature and then to
incorrect conclusions on chemical composition and evolution stage
since the expected values alterations of abundances in evolution
are small, they exceed possible systematic errors insignificantly.
Therefore we determined the effective temperature by the
traditional method, from the condition of independence of neutral
iron abundance upon excitation potentials of the corresponding
lines (Fig. 1a). The surface gravity was chosen from the
ionization balance for neutral and ionized atoms of iron, and the
microturbulent velocity from the condition of independence of
iron abundance and line intensity (Fig. 1b).
When determining the model atmospheric parameters and calculating
chemical composition, it is important to restrict the study to
lines of low and moderate intensities, with equivalent widths
Å, since the stationary plane-parallel
atmosphere approximation can be inadequate for the presentation
of the strongest spectral features. Besides, some of the strong
absorptions can be distorted by the circumstellar envelope and,
if the spectral resolution is not sufficient, the intensity of
the envelope components will be included in the intensity of the
components formed in the atmosphere.
The equivalent widths, oscillator strengths (gf) of the spectral lines used to determine the model parameters, and the computed abundances of elements are presented in Table "K413-LIN'' available upon request by , and via ftp at the FTP-directory: lnf1.sai.msu.ru/pub/PEOPLE/samus/ K413-LIN (or ftp.sao.ru/pub2/sslab/K413-LIN).
The effective temperature was thus determined from their dependence of the iron abundance on the excitation potential
of the lower levels of the lines used. An additional criterion
for the method's reliability is the absence of a similar
dependence for other chemical elements, also represented in the
spectra by numerous lines (for example, SiI, CaI, ScII,
TiI, TiII, CrI, CrII, NiI). Besides, in the case of
reliable determination of the microturbulent velocity
,
individual abundances are independent of the equivalent
widths of the lines used in computations. The abundances of
silicon, titanium, and chromium, determined from lines of
neutral atoms as well as from lines of ions, agree within the
uncertainties of the method. This evidence for the correctness
of our determination of the atmospheric gravity comes from the
ionization balance for atoms of iron. Generally, the
self-consistency of the parameters shows that the homogeneous
model atmospheres applied here are usable for computations of
weak lines in the LTE approximation. The resulting parameters of
the model atmospheres are,
=4800K, logg=0.7,
=3.5 km s-1.
![]() |
Figure 1: a) Iron abundance FeI (circles) and FeII (crosses) calculated for K413 with model parameters from Table 3 using lines with different excitation potencials EP of a low lewel. |
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![]() |
Figure 1: b) The same as a function of the equivalent width W. |
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The characteristic uncertainty of model parameter determinations, on
average for a star with temperature about 5000K, is
K,
log
0.3dex,
0.5 km s-1.
The corresponding abundances shifts
caused by uncertainties in atmospheric parameters are tabulated in
Table 2.
Mainly, they do not exceed 0.15dex. The main exception is an oxygen
abundance determination error (as large as 0.16dex) due to its
sensitivity to temperature and gravity.
Besides, the barium abundance, based on relatively strong lines,
,
could be disturbed by a large systematic error due to the sensitivity of
its abundance to the microturbulent velocity value. We have to keep
in mind, first, that the atmospheric parameters are not independent
and such estimations of uncertainties are not absolutely valid and,
second, that their influence on the relative abundances [X/Fe] is
essentially reduced.
The dominant source of inaccuracy of chemical abundances is associated
with inaccuracy of equivalent widths. Most of the lines used for
computations of the chemical composition have equivalent widths Wbelow 100mÅ, resulting in very high demands on the quality of the
observing material, since the accuracy of W values for weak lines, at
a given spectral resolution, mainly depends upon the signal-to-noise
ratio in the spectrum. The line-to-line scatter in the abundances
(Table 3) is mainly contributed by
uncertainities in measured equivalent widths W and in the
gf-values. The scatter of element abundances derived from numerous
(n > 10) lines is not so high, the rms deviation,
,
generally does not exceed 0.25dex
(cf. Table 3). It should be noted here that for elements having
less than five lines the errors calculated are not a real estimation of
an internal abundance scatter.
On the whole, we can expect abundances of elements with a sufficiently large number of lines (Si, Ca, Sc, Ti, Fe-peak) and gf-values of good quality, to be accurate within 0.20 to 0.25dex and for other elements the uncertainty could be above 0.25-0.30dex.
All computations in this study used the WIDTH9 code developed by Tsymbal (1996) for LTE. Corrections for hyperfine structure and for isotope shift causing broadening of Ni I, Mn I, and Ba II lines were not taken into account.
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|||
Species |
![]() |
![]() |
![]() |
CI | 0.11 | 0.15 | 0.01 |
OI | 0.16 | 0.16 | 0.02 |
NaI | 0.04 | 0.03 | 0.03 |
MgI | 0.05 | 0.03 | 0.13 |
AlI | 0.02 | 0.02 | 0.02 |
SiI | 0.02 | 0.01 | 0.02 |
SiII | 0.13 | 0.17 | 0.03 |
SI | 0.09 | 0.13 | 0.01 |
CaI | 0.06 | 0.04 | 0.12 |
ScII | 0.01 | 0.13 | 0.07 |
TiI | 0.10 | 0.05 | 0.04 |
TiII | 0.03 | 0.11 | 0.12 |
VI | 0.09 | 0.05 | 0.00 |
VII | 0.04 | 0.12 | 0.00 |
CrI | 0.08 | 0.05 | 0.07 |
CrII | 0.07 | 0.12 | 0.01 |
MnI | 0.09 | 0.05 | 0.02 |
FeI | 0.07 | 0.04 | 0.09 |
FeII | 0.04 | 0.13 | 0.05 |
NiI | 0.08 | 0.03 | 0.04 |
CuI | 0.09 | 0.03 | 0.02 |
ZnI | 0.05 | 0.04 | 0.08 |
YII | 0.01 | 0.11 | 0.04 |
ZrI | 0.12 | 0.06 | 0.00 |
ZrII | 0.01 | 0.12 | 0.01 |
BaII | 0.02 | 0.12 | 0.30 |
LaII | 0.01 | 0.12 | 0.00 |
CeII | 0.01 | 0.11 | 0.01 |
NdII | 0.02 | 0.12 | 0.01 |
PrII | 0.01 | 0.10 | 0.00 |
EuII | 0.00 | 0.12 | 0.01 |
The spectrum of K 413 exhibits several peculiar features,
in particular, a complex absorption and emission profile of
the H
line (cf. Fig. 2).
One can assume that the observed H
profile is a
superposition of the photospheric absorption and the
line formed in higher layers. Figure 2 shows the
profiles of the H
line for two observing moments in
comparison with a photospheric (theoretical) profile.
Such profiles were first found in the spectra of four globular
cluster stars by Cohen (1976) using photographic echelle
spectra recorded with a two-stage image tube. Then Mallia &
Pagel (1978) raised the number of H
-emission
stars to 12 with the RGO spectrograph and a photon counter.
![]() |
Figure 2:
Same as in Fig. 1 but for the wavelength region
near the H![]() |
Open with DEXTER |
Recording spectra with a diode array at a three times poorer
resolution compared to our resolution, Cacciari & Freeman
(1983) surveyed spectra of bright stars in 12 globular clusters. One third of the stars with luminosities
higher than solar by three orders of magnitude showed
indications of H
emission. The effect does not correlate
with metallicity of the corresponding cluster. It follows from
data in Table3 of the cited paper that the velocity difference
between the absorption core and the emission peaks does not
significantly change from star to star, being about 50 km s-1. The
mass loss rate was determined in the model of stellar wind, like
in the studies by Cohen (1976) and by Mallia & Pagel
(1978). Our Table 1 shows that the velocity
difference between the core and the emission peaks for K 413 is
in agreement with the general pattern of the Cacciari & Freeman
(1983) survey.
Under the assumption of emission formed in an extended circumstellar
envelope, the observed emission profiles can be obtained both for
accretion on the star and for loss of matter. Indications of
neither interstellar hydrogen (Knapp et al. 1973) nor
dust (the results of the IRAS survey analyzed by Knapp et al.
1995) were found in globular clusters, thus the
hypothesis of wind outflow of matter, with subsequent loss of
this matter from the cluster as a whole, seems natural. The
difference between the H
emission of K 413 and the
classical PCygni profile can be due to a thinner circumstellar
envelope and a lower mass loss rate. Low optical thickness
follows also from the fact that no contribution from the envelope
is observed for other strong lines of the photospheric spectrum
(with the exception of the sodium resonance doublet). Mass loss
rate estimates from Cacciari & Freeman (1983), Cohen
(1976), Mallia & Pagel (1978) were essential
for the problem of consistency of evolutionary computations of
the RGB, HB, and AGB stages. It was noticed already in the first
observations that the intensity ratio of the short-wave and
long-wave residuals of the emission component differed from star
to star and varied in time for each star. In the circumstellar
envelope model, this means matter infall as well as outflow.
However, the short time scale of such variations (several days)
caused doubt on the hypothesis of an extended circumstellar
envelope.
Sun | K 413 |
![]() |
|||||
-1.38, 4800 K, 0.7, 3.5 km s-1 | -1.48, 4790 K, 1.3, 1.2 km s-1 | ||||||
Species |
![]() |
X |
![]() |
n | ![]() |
![]() |
![]() |
C | 8.55 | CI | 8.35 | 9 | 0.44 |
![]() |
+1.10 |
O | 8.87 | OI | 9.66 | 3 | 0.15 | +2.17 | +0.91 |
Na | 6.33 | NaI | 5.39 | 5 | 0.38 | +0.44 | -0.14 |
Mg | 7.58 | MgI | 6.81 | 4 | 0.04 | +0.61 | +0.26 |
Al | 6.47 | AlI | 6.26 | 2 | +1.17 | +0.07 | |
Si | 7.55 | SiI | 6.68 | 19 | 0.26 | +0.51 | +0.28 |
SiII | 6.72 | 1 | +0.55 | +0.33 | |||
S | 7.21 | SI | 7.01 | 4 | 0.24 | +1.18 | |
Ca | 6.36 | CaI | 5.22 | 21 | 0.18 | +0.24 | +0.30 |
Sc | 3.17 | ScII | 1.62 | 11 | 0.13 | -0.17 | -0.22 |
Ti | 5.02 | TiI | 3.74 | 27 | 0.27 | +0.10 | -0.02 |
TiII | 3.73 | 15 | 0.16 | +0.09 | +0.15 | ||
V | 4.00 | VI | 2.76 | 9 | 0.28 | +0.14 | -0.36 |
VII | 2.71 | 4 | 0.24 | +0.09 | +0.15 | ||
Cr | 5.67 | CrI | 4.28 | 15 | 0.24 | -0.01 | -0.24 |
CrII | 4.29 | 7 | 0.17 | 0.00 | -0.04 | ||
Mn | 5.39 | MnI | 4.02 | 5 | 0.37 | +0.01 | -0.13 |
Fe | 7.50 | FeI | 6.14 | 180 | 0.24 | +0.02 | +0.01 |
FeII | 6.09 | 22 | 0.23 | -0.03 | -0.02 | ||
Ni | 6.25 | NiI | 4.80 | 26 | 0.18 | -0.07 | -0.12 |
Cu | 4.21 | CuI | 2.41 | 3 | 0.22 | -0.42 | -0.54 |
Zn | 4.60 | ZnI | 3.36 | 4 | 0.17 | +0.14 | -0.01 |
Y | 2.24 | YII | 0.55 | 7 | 0.26 | -0.31 | -0.57 |
Zr | 2.60 | ZrI | 1.51 | 3 | 0.37 | +0.29 | +0.02 |
ZrII | 1.39 | 3 | 0.26 | +0.17 | -0.26 | ||
Ba | 2.13 | BaII | 0.63 | 3 | 0.19 | -0.12 | +0.04 |
La | 1.22 | LaII | -0.23 | 3 | 0.29 | -0.07 | +0.20 |
Ce | 1.55 | CeII | 0.22 | 5 | 0.20 | +0.05 | -0.22 |
Nd | 1.50 | NdII | -0.09 | 9 | 0.26 | -0.21 | -0.11 |
Pr | 0.71 | PrII | -0.46 | 2 | +0.21 | +0.17 | |
Eu | 0.51 | EuII | -0.39 | 3 | 0.19 | +0.48 | +0.05 |
![]() ![]() ![]() |
Other authors preferred the hypothesis that the Hemission was formed in the lower chromosphere of cool luminous stars.
Formation of the chromospheric absorption H
component is
photoionization-controlled; this conclusion, first formulated for
the Sun, was long ago extended to most types of cool stars. An
exception is represented by dMe dwarfs where photoionization
corrections to the source function are low compared to impact
corrections (Cram & Mullan 1979). Consequently, for most
types of cool stars, the source function is insensitive to the
structure of the chromosphere and is determined only by effective
temperature and gravity. Cram & Mullan (1985) pointed
out that observed H
equivalent widths in the spectra of
cool stars (
K) were considerably higher than
theoretical photospheric ones, whereas addition of a simple model
chromosphere (two segments, with a linear change of temperature in
each of them) made it possible to eliminate the discrepancy. They
concluded that each cool star possessed a chromosphere with a
significant optical depth in H
.
They also confirmed the
old idea of nonthermal broadening of absorption profiles by the
velocity field in the chromosphere. For supergiants, the pattern
is more complicated than the cited models: somewhere for
log
,
increasing outflow rate must cause the
situation when the hydrostatic approximation is not valid for the
formation region of the H
core. Our observations confirm
both conclusions: the central absorption width of the observed
H
profile significantly exceeds the theoretical
photospheric value, and the "residual'' absorption exhibits a
displacement of the core towards shorter wavelengths. The
"blue'' displacement of the absorption core evidence
outflow; more exactly, the most peripheral of the core-forming
layers move outward with respect to the layers forming the
absorption profile at half intensity. The excess intensity of the
short-wavelength residual of the chromospheric emission peak over
the long-wavelength one can be interpreted either as infall of
absorption-forming layers with respect to emission-forming layers
or as expansion of these layers, with deceleration (relative to
the emission region). Note that the non-thermal broadening in the
emission-forming region exceeds that for the external
chromosphere.
Dupree et al. (1984) demonstrated that Hemission could be formed in dense, hot chromospheres. The high
sensitivity of the profile to the chromospheric temperature
permits us to conclude that the presence of emission wings is not
direct proof of matter outflow. Conclusive evidence can be
asymmetry of the central absorption or different intensities of
residual emission peaks. Smith & Dupree (1988) studied
spectra of 52 halo giants (echelle spectrograph with Reticon)
and found emission wings of the H
line for 10 stars with
.
Profile widths, their variability and different
ratios of the short-wave component to the long-wave can be
explained by differential motions in the chromospheric region
with relatively high density and
K to
8000K, where the emission is formed. Adding data on
chromospheric components of the resonance MgII and CaII
lines made it possible to finally conclude in favor of the
chromospheric origin of the H
emission (Dupree et al.
1990).
Besides progress in model description, the development of the
chromospheric hypothesis is due to improvement of observational
possibilities. With better spectroscopic resolution and
signal-to-noise ratio, the fraction of luminous stars with
revealed H
emission increases: it was 30% in Cacciari
& Freeman (1983), 60% in Gratton et al.
(1984), 72% in Bates et al. (1993). The
displacement of the H
absorption core correlates with
the star's luminosity and with the core displacements for the
resonance sodium doublet lines (Kemp & Bates 1995).
As it can be seen in Fig. 2, the H
profile
in the K 413 spectrum is time variable,
which also confirms the chromospheric origin of its emission.
The second characteristic feature of the spectrum of K 413 is
the presence of several absorptions we identify with diffuse
interstellar bands (DIBs). The technique we use to isolate such
bands is the same as in Klochkova et al. (2000). The
equivalent widths of the most easily detectable bands, 6203, 6379, 6613 Å Å, are respectively W=24, 20, and
66m Å.
The best-known band,
5780 (Fig. 3),
is blended with the star's lines and so it
is impossible to measure its parameters separately. Thus,
absorption bands, most probably formed in the stellar vicinity, have
been revealed in the spectrum of an evolved star, a
globular cluster member. Earlier, such spectral features were
found for several evolved stars: Cohen & Jones (1987)
discovered diffuse interstellar bands in the spectrum of the
heavily reddened nucleus of the planetary nebula WC11; Le Bertre
& Lequeux (1993) studied a sample of stars featuring mass loss
and revealed diffuse interstellar bands in the spectra of
IRC+10420 and ACHer. In the cause of the present program, we
also found diffuse interstellar absorptions in the spectra of
several cool post-AGB objects (cf. Klochkova 2000 and
references therein).
![]() |
Figure 3:
Same as in Fig. 1 but for the wavelength region
with the strongest DIB ![]() |
Open with DEXTER |
From the cluster membership of K 413, taking into account the
mean color excess,
E(B-V) = 0.19, and the distance
modulus,
(Harris 1996),
we obtain the star's absolute magnitude,
.
The real reddening value for K 413 can be much higher because
of the additional circumstellar reddening revealing itself in the
absorption bands (like the 5780 Å band in Fig. 3).
In such a case, the star's luminosity can be considerably higher.
Revealing absorption bands in the spectrum of a globular-cluster star is an unexpected result: as stated above, searches for dust in the regions of globular clusters, making use of the IRAS data, gave negative results (Lynch & Rossano 1990).
Until now, estimates of metallicity for the cluster M12 were
based only on photometric data or on low-resolution spectra.
From observations of the Ca II triplet (8498, 8542, 8662 Å),
Da Costa & Armandroff (1995) derived
and attributed M12 to the intermediate subsystem of globular
clusters. However, Rutledge et al. (1997), also from
low-resolution intensities of the Ca II triplet, found a higher
metal abundance,
.
Carretta & Gratton
(1997), from high-resolution spectroscopic
observations of several stars in selected globular clusters,
suggested a new uniform metallicity scale for a large sample of
globular clusters.
For M12, they obtained
.
Chemical element abundances,
,
averaged over the measured lines, are collected in
Table 3. Its second column reproduces, from Grevesse
et al. (1996), the corresponding data for the solar
atmosphere we use to determine the relative abundances,
For comparison, Table 3 presents the chemical
composition of
,
a luminous star at a high
galactic latitude. Klochkova et al. (2001a), from a variety
of its characteristics, concluded that the high-latitude supergiant
,
having
,
probably belonged to the early post-AGB evolution stage.
![]() |
Figure 4:
Same as in Fig. 1 but for the wavelength region
of the spectrum of K 413 containing OI oxygen
lines near 6156 Å. The dashed line shows the position of the
continuum. The thin line shows the theoretical spectrum calculated
with the model parameters and chemical compostion from the table, but the
value
![]() |
Open with DEXTER |
The metallicity value derived by us for K 413,
,
is in a good agreement with the results
of Carretta & Gratton (1997). The mean abundance of
the iron-peak metals (vanadium, chromium, manganese, nickel) does
not differ considerably from that of iron: for K 413,
.
Copper, being an iron-peak element not
universally recognized (Timmes et al. 1995), shows a slight
deficiency,
;
this value has been reliably
established from 3 lines. Sneden et al. (1991)
concluded that similar copper underabundance was observed for
halo stars in globular clusters as well as in the galactic field.
The most interesting feature in the distribution of abundances
for K 413 is the large oxygen excess (
)
revealed from 3 confidently measured OI lines around 6156 Å (see Table 4). Unfortunately, the lines of the IR
oxygen triplet
7773 Å in the available spectrum of
K 413 are at the very edge of the echelle frame, so we did not
apply these lines either for oxygen abundance or luminosity
determinations. To illustrate the determination of [O/Fe]
ratio, we present Fig. 4 showing a comparison between
the observed spectrum and the theoretical one calculated for two values
[O/Fe]=+2.2 and [O/Fe]=+0.0 computed with the STARSP
code (Tsymbal 1996). This oxygen overabundance,
,
is much higher than the value observed in
the atmospheres of halo unevolved stars (Timmes et al.
1995), so we conclude the presence of dredge-up of oxygen produced in
helium burning.
![]() |
EP, eV | lggf | W |
![]() |
OI | ||||
6155.98 | 10.74 | -0.66 | 12 | 9.83 |
6156.77 | 10.74 | -0.44 | 11 | 9.56 |
6158.18 | 10.74 | -0.29 | 15 | 9.60 |
7771.94 | 9.15 | 0.33 | 105 | 9.01 |
7775.39 | 9.15 | -0.03 | 99 | 9.30 |
The carbon abundance,
,
has been determined
from nine CI lines in the spectrum of K 413. This value has
a large uncertainty because of low intensity of the CI lines.
Besides, the most reliable lines around 7110 Å are blended with
telluric lines.
The nitrogen lines, NI 7423, 7442, 7468 Å Å, are strong
enough, but since the wavelength region with these lines was
recorded at the edge of our echelle frame, where line positions
could be measured with poor accuracy, we do not present the
nitrogen abundance because of its low reliability.
From four lines, we derive a normal zinc abundance, relative to iron,
for K 413:
.
Just the abundance of zinc, not
changing in the course of the stellar nucleosynthesis in the
interiors of low- and intermediate-mass stars, varies over a wide
metallicity range by the same amount as the iron abundance (Sneden et
al. 1991). Thus the conclusion on zinc abundance is
independent of the scale used (a relative or absolute one). This fact,
together with the practically normal value,
,
leads to the conclusion of no selective separation of chemical elements
in the probable circumstellar envelope of K 413. This conclusion is
also supported by the good agreement between the metallicity of
K 413 and that for the cluster as a whole.
Now consider the abundances of light metals in the atmosphere of
K 413. From the group of odd -process elements, we have
obtained data for sodium and aluminum. Sodium is slightly enhanced,
;
Denissenkov & Ivanov (1987) and
Denissenkov (1989) showed that such an excess could be due
to synthesis of sodium in the process of hydrogen burning. However,
the obtained sodium excess is probably due, in part, to no treatment
of deviations from the local thermodynamic equilibrium (LTE).
Mashonkina et al. (2000) demonstrated that, for a luminous
star with
K, accounting for
superionization of sodium atoms leads to corrections of about
-0.2dex to sodium abundance, compared to the LTE approximation,
for the subordinate sodium lines used by us. Consequently, a
considerable part of the sodium excess revealed by us is probably of
methodogical, not evolutionary, origin.
The high abundance of aluminum derived by us from two lines is not a rare case for stars in globular clusters with moderate metal deficiency, as it follows from Shetrone (1996). However, K 413 is peculiar in the way that its aluminum excess is accompanied by oxygen excess, whereas, according to Shetrone (1996), an Al-O anticorrelation is observed for globular cluster giants. We explain this anomaly of the chemical composition of K 413 by its more advanced evolutionary stage.
Of even -process elements, we determined abundances for
Mg, Si, S, Ca, and Ti. Note that titanium (like scandium) is an
intermediate element, sometimes considered the heaviest
-process element and sometimes, the lightest iron-peak
metal (Wheeler et al. 1989). In the case of K 413, the
abundances of Sc and Ti correspond to that of iron.
The excess silicon abundance,
,
has been
reliably determined from two ionization states. Similar relative
silicon abundances,
,
are also observed in
the atmospheres of unevolved stars with metallicities close to
that of K 413 (Timmes et al. 1995). Consequently, from our
data, we can conclude a normal silicon abundance for K 413.
However, the silicon excess in the atmosphere of
,
a luminous star in the galactic field (Table 3),
is somewhat lower.
Klochkova et al. (2001a) have already noted that sulphur
abundances in the atmospheres of evolved stars seem to be a
special problem. Earlier, Bond & Luck (1987) revealed a
large sulphur excess,
,
in the atmosphere
of the low-metallicity post-AGB star HD46703 and explained it
with probable synthesis of sulphur by addition of
particles to
nuclei. However, Klochkova (1995)
found an excess abundance of sulphur in the atmosphere of the
normal massive supergiant
Per. First, it has no
circumstellar envelope and thus there are no condensation
processes; second, we cannot expect any manifestations of
synthesis of sulphur at such an early evolutionary stage.
Sulphur excess seems to be a persistent feature in the pattern
of the chemical composition of evolved stars (besides the results in
Bond & Luck 1987, cf., for instance, Klochkova
1995; Van Winckel et al. 1996; Klochkova et al.
1999, 2000).
For these reasons, we still rather tend to explain the
observed excess of sulphur in the atmosphere of K 413 by
methodotogical reasons.
It should be noted that the abundance of sulphur, relative to iron, is also enhanced in the atmospheres of unevolved stars (Timmes et al. 1995 and references therein). Recently were published also new results concerning sulphur abundances. Using the near-infrared S I doublet, Israelian & Rebolo (2001) derived enhanced sulphur abundances for some unevolved metal-poor stars. These authors concluded that an the increasing of [S/Fe] ratio for most metal-deficient stars could be explained through the increasing role of of hypernovae in the earliest evolution of the Galaxy.
A very important characteristic of the chemical composition of
evolved stars is the abundance of heavy elements, formed in the
synthesis due to the slow process of neutron capture. It is
usually accepted (Schwarzschild & Härm 1965;
Lattanzio & Forestini 1999; Blöcker 2001)
that an excess of heavy metals can be observed in the
atmospheres of post-AGB stars, due to neutron capture, mixing,
and dredge-up of matter processed in the interiors to the
surface. The data in Table 3 rather show depleted
abundances of these elements in the atmosphere of K 413. In
particular, the abundance of barium, confidently determined from
three lines, is
,
in agreement with the
general trend of this ratio for halo stars (McWilliam 1997).
The relative abundances of lighter s-process metals, Y and Zr,
have been determined less accurately because of the limited number
of lines available, but on average in this case, too, their
relative abundance,
,
appears to be close
to the solar value.
We noted earlier (Klochkova et al. 2001b) that deficiency of s-process elements in the atmospheres of supergiants, also at the post-AGB stage, was observed much more often than their excess. Both physical and methodical reasons for the deficiency were frequently discussed (see, for example, Luck & Bond 1989). In our opinion, the lack of heavy elements' dredge-up manifestations observed for most supergiants is real and not due to systematic errors of model atmosphere analysis for spectra of supergiants. Most probably, presence or absence of excess for s-process elements is somehow related to such fundamental parameters of the star as its initial mass and mass loss rate, determining the evolution of an individual star.
As for heavier elements, with
(La, Ce, Nd, Pr),
we can generally speak about complete agreement with the solar
relative abundance of these elements:
.
Such
a pattern agrees with a behaviour of lanthanides in halo and in
globular cluster stars (Travaglio et al. 1999; Wheeler et al.
1989). At the moderate metal deficiency, these elements
represent products of both slow and rapid neutron captures.
A slight excess of europium (created only in the r-process),
,
is typical of stars in low-metallicity
globular clusters (Shetrone 1996).
The value log(Ba/Eu) is a traditional contribution indicator
for r- and s-processes for the synthesis of these metals.
The value
for K 413 is in agreement
with the general pattern for the halo (Spite 1992;
Timmes et al. 1995).
Comparing the data in Table 3, we can conclude that the
main parameters and details of the chemical composition of K 413 are
close to the corresponding parameters for
.
Taken together, the features of the chemical composition (excess of
oxygen and of
-process elements), combined with the star's
position in the color-magnitude diagram, make us assume the post-AGB
evolutionary stage for K 413.
Such post-AGB objects are frequently met in globular clusters. For example, Huges and Wallerstein (1997) attributed the star N13 (with the absolute magnitude MV = -2.06) in the cluster M15 to this stage of evolution. It appears from the color-magnitude diagram (Brocato et al. 1996) that M12 has other stars with V < 13 and (B-V) < 1.
From the main peculiarities of its chemical composition (excess
of O, excess of light metals: Na, Mg, Al; no excess of s-process
elements), K 413 is a hotter analog of the "CO-normal'' group
of stars, Nos.74, 91, and 256 in the cluster
Cen, the
latter's chemical composition studied by Francois et al.
(1988). Note, however, that, for the stars 74, 91, and
256, the relative abundances of the s-process elements,
[heavy/Fe], are depleted by the factor of 3 to 4 compared to the
solar value, whereas for K 413, the value of [heavy/Fe] is
closer to that for the Sun. By the presence of emission in the
wings of the H
line, K 413 is similar to the star 65 in
Cen.
At the same time, the chemical composition of K 413 differs
drastically from that of the highly evolved star ROA24
(Ferenbach's star) in the cluster Cen.
Gonzalez & Wallerstein (1992) determined the detailed
chemical composition of ROA24 from high-resolution spectra.
The large excesses of CNO elements and s-process metals in its
atmosphere, along with its high luminosity, evidence for the AGB
(or post-AGB) evolutionary stage. It should be noted that the
significantly different effective temperatures of K 413 and
ROA24 cannot be a key issue for explanation of the principal
differences between the chemical compositions of the two stars.
The cluster
Cen contains a sample of luminous stars
with effective temperatures around 4100-4500K, both with
barium excesses and barium deficiencies (cf. Fig. 3 in Gonzalez
& Wallerstein 1992). Later, Gonzalez & Wallerstein
(1994) came to the conclusion that the main parameter
determining the dredge-up of the synthesis products to a star's
atmosphere at the post-AGB stage was its luminosity.
Their conclusion is confirmed by the fact that ROA24 is the
most luminous star in the Galaxy's system of globular clusters.
Detailed information of the velocity pattern in the atmosphere is needed to clarify the object's evolutionary status. The results of our radial velocity measurements for K 413 are presented in Table 1. To improve accuracy of velocity determinations, we compared the observed spectrum with the corresponding synthetic one and selected unblended lines. The positional zero point for each spectrogram was determined in the standard way, relative to ionospheric emissions of the night sky and to the absorption telluric spectrum, recorded with the object's spectrum.
The typical uncertainty of the average was
km s-1, for the number of lines exceeding 400
and the measurement uncertainty for a single line
km s-1.
The mean velocity from lines of metals,
km s-1 for the spectrum s27605 and
km s-1 for the second one, can be considered
the systemic velocity; it is in good agreement with the radial
velocity value,
km s-1, found by Harris et al.
(1983) for the whole cluster from low-resolution spectra.
The agreement between radial velocities of K 413 obtained in this
study and by Harris et al. (1983) more than 20 years earlier
permits us to conclude that the star's radial velocity is constant.
Thus, we can reject the suspicion of possible binarity of K 413,
the star occupying a peculiar position in the color-magnitude diagram.
As noted above, the H
line has a complex emission and
absorption profile with two emission peaks, at the velocities
and
km s-1 for the spectrum s27605, differing from the systemic velocity
correspondingly by
and +58.3 km s-1.
It follows from the data in Table 1 that the absorption
core of the H
line has a velocity differing from the
mean metallic-line radial velocity of the star by more than
one rms error. This displacement of the absorption
component towards shorter wavelengths is also clearly seen in
Fig. 2, where the position of the core of the theoretical
H
profile corresponds to the star's velocity derived
from the absorption lines of metals.
The mean radial velocity values measured from several diffuse
bands we identified in the spectrum of K 413 are collected in
Table 1. It is natural to assume that these absorptions
are formed in the circumstellar region: their interstellar origin
is improbable because of the cluster's high position in the
Galaxy (
). The difference between
the radial velocities measured in the spectrum of K 413 from lines
formed in the photosphere and from those formed in the circumstellar
envelope gives a negative value for the velocity of the envelope's
motion with respect to the central star,
km s-1,
and thus we have reason to consider infall of the circumstellar
matter onto the star.
It was already noted above that, according to Smith & Dupree
(1988), the excess intensity of the short-wavelength emission
component of the H
line over the long-wavelength one
(such an intensity ratio of the emission peaks is characteristic
of K 413) is evidence is for decelerating matter outflow in the
star's chromosphere. But we see the opposite behavior in two H
profiles obtained for two observing moments.
The lines of the resonance sodium doublet, NaD1, 2, are
asymmetric in the spectrum of K 413, their cores shifted to
longer wavelengths. Measurements show that both lines are blends,
poorly resolved at our spectral resolution, and each of them
consists of two components (see Table 1). One of the
components has the velocity coinciding with that of the star, and
the position of the second component corresponds to the velocity
measured from diffuse bands. Peterson (1981) was the
first to notice asymmetry of the sodium doublet lines in the
spectra of globular-cluster giants. A detailed study of
displacements of the sodium doublet lines with respect to stellar
velocities, based upon high-resolution spectra, was performed by
Bates et al. (1993) for the clusters M22 and
Cen. For a sample of program stars, these authors
revealed displacements of the doublet lines approximately within
the -10 to +2 km s-1 range, interpreted as a manifestation of
matter outflow (or infall, in the cases of positive
displacements) in stellar envelopes. In particular, analyzing
their data, Bates et al. (1993) concluded that
displacements of the Na doublet lines were observed for luminous
stars,
.
This result can be
considered additional evidence in favor of our assumption of
the higher luminosity of K 413 compared to the value derived from
its cluster membership.
From CCD spectra obtained with the echelle spectrometer of the 6
m telescope, using model atmospheres, we determined the
fundamental parameters
=4800K, logg=0.7
and the detailed chemical composition for the star K413
in the globular cluster M12. The measured radial velocity of the star for two observing moments
and
km s-1,
agrees very well with the mean velocity of M12, confirming cluster
membership of the star.
The most important feature in the optical spectrum of K413 is
the presence of emission components in the wings of the Hline, with two variable emission peaks.
The spectrum of K413 contains absorption bands, with positions
making it possible to identify them with the so-called diffuse
interstellar bands. The difference between the radial velocity
corresponding to these bands and the star's velocity is
km s-1. The sodium doublet lines, Na D1, 2,
are blends, with one of their components having the same displacement
towards longer wavelengths.
The obtained value
is the first
determination of metallicity for the cluster M12 from
high-resolution spectra. This metallicity value agrees well with
the published mean metallicity of M12, derived from photometry
and low-resolution spectroscopy.
The principal anomaly of the chemical composition of the star's
atmosphere is a large oxygen excess,
.
The s-process metals are depleted relative to metallicity: for Y
and Zr,
,
and for barium,
.
The abundances of heavier elements: La,
Ce, Nd, and Pr, do not differ from solar with respect to
iron:
.
The europium excess,
,
is characteristic of stars in low-metallicity
globular clusters.
The star's position in the color-magnitude diagram of the cluster M12 and its chemical composition permit us to suppose that K413 is in the post-AGB stage of evolution.
Acknowledgements
The authors are grateful to N. S. Tavolganskaya for her great help in spectral reductions.
We are much indebted to our referee for critical reading of the manuscript and for his valuable notes and comments.
One of the authors (V.G.K.) wishes to thank the Russian Foundation for Basic Research for financial support of the spectroscopic investigations, with the 6-meter telescope, of objects evolving from AGB stars to planetary nebulae (project 99-02-18339) and the Federal "Astronomy'' Program (project 1.4.1.1.). The research described in this publication was made possible in part by Award No.RP1-2264 of the the U.S. Civil Research & Development Foundation for the Independent States of the Former Soviet Union (CRDF).
N.N.S. would like to thank the council of the program of support for Russia's leading scientific schools, for financial support (grant 00-15-96627).
This study made use of the SIMBAD astronomical data base, of the CDS bibliographic data base, and of the VALD atomic data base.