A&A 462, 711-730 (2007)
DOI: 10.1051/0004-6361:20065785
D. A. García-Hernández1 - P. García-Lario1,2 - B. Plez3 - A. Manchado4,5 - F. D'Antona6 - J. Lub7 - H. Habing7
1 - ISO Data Centre, Research and
Scientific Support Department of ESA. European Space Astronomy Centre (ESAC), Villafranca del
Castillo, PO Box 50727, 28080 Madrid, Spain
2 -
Herschel Science Centre, Research and Scientific Support Department of ESA. European Space Astronomy Centre (ESAC), Villafranca del
Castillo, PO Box 50727, 28080 Madrid, Spain
3 -
GRAAL, CNRS UMR 5024, Université de Montpellier 2, 34095 Montpellier
Cedex 5, France
4 - Instituto de Astrofísica de Canarias, La Laguna 38200, Tenerife,
Spain
5 - Consejo Superior de Investigaciones Científicas (CSIC), Spain
6 - Osservatorio Astronomico di Roma, via Frascati 33, 00040 MontePorzio
Catone, Italy
7 - Sterrewacht Leiden, Niels Bohrweg 2, 2333 RA Leiden, The Netherlands
Received 8 June 2006 / Accepted 25 August 2006
Abstract
Lithium and zirconium abundances (the latter taken as representative
of s-process enrichment) are determined for a large sample of massive
Galactic O-rich AGB stars, for which high-resolution optical spectroscopy has
been obtained (
). This was done by computing synthetic
spectra based on classical hydrostatic model atmospheres for cool stars and using
extensive line lists. The results are discussed in the framework of
"hot bottom burning'' (HBB) and nucleosynthesis models. The complete sample is
studied for various observational properties such as the position of the
stars in the IRAS two-colour diagram ([12] - [25] vs. [25] - [60]),
Galactic distribution, expansion velocity (derived from the OH maser emission),
and period of variability (when available). We conclude that a considerable
fraction of these sources are actually massive AGB stars (M>3-4
)
experiencing HBB, as deduced from the strong Li overabundances
we found. A comparison of our results with similar studies carried out in the past
for the Magellanic Clouds (MCs) reveals that, in contrast to MC AGB stars,
our Galactic sample does not show any indication of s-process element
enrichment. The differences observed are explained as a consequence of
metallicity effects. Finally, we discuss the results obtained in the framework
of stellar evolution by comparing our results with the data available in the
literature for Galactic post-AGB stars and PNe.
Key words: stars: AGB and post-AGB - stars: abundances - stars: evolution - nuclear reactions, nucleosynthesis, abundances - stars: atmospheres - stars: late-type
Another important characteristic of AGB stars is the presence of neutron-rich
elements (s-elements such as Sr, Y, Zr, Ba, La, Nd, Tc, etc.) in their
atmospheres. These species are formed by slow-neutron captures in the
intershell region. After a dredge-up episode, according to theoretical models,
protons can be partially mixed in the 12C- and 4He-rich region
between the hydrogen- and helium-burning shells (e.g. Straniero et al. 1995,
1997; Mowlavi 2002; Lattanzio 2003) and can react with 12C to form
13C during the interpulse phase. The 13C(,n)16O reaction
releases neutrons, which are captured by iron nuclei and other heavy elements,
forming s-elements that can later be dredged up to the stellar surface in the
next thermal pulse (Straniero et al. 1995, 1997; Busso et al. 2001).
Another neutron source, 22Ne, needs to be considered. In this case the
activation takes place during the convective thermal pulses (e.g. Straniero et al. 1995, 2000; Vaglio et al. 1999; Gallino et al. 2000). 22Ne can be
formed from 14N, and 14N is formed by the CNO cycle in the hydrogen
shell. The 22Ne(,n)25Mg neutron source requires higher
temperatures and typically occurs at higher neutron densities than the
13C(
,n)16O reaction. Thus, a different s-element pattern is
expected depending on the dominant neutron source. According to the
most recent models, the 13C(
,n)16O reaction is the preferred
neutron source for masses around 1-3
,
while for more massive stars
(i.e.
)
neutrons are thought to be mainly released
through the 22Ne(
,n)25Mg reaction (see, for example, Busso et al. 1999 and Lattanzio & Lugaro 2005, for a recent review). In the
literature, there is strong evidence that most Galactic AGB stars enriched in
s-process elements have masses around 1-3
,
where the
13C(
,n)16O reaction is the neutron donor (e.g. Lambert et al. 1995; Abia et al. 2001). Unfortunately, a confrontation of the predictions made
by these models with observations of more massive AGB stars in our Galaxy is
not yet available.
In the case of the more massive O-rich AGB stars (
), the
convective envelope can penetrate the H-burning shell thereby activating so-called
"hot bottom burning'' (hereafter, HBB), which takes place when the temperature
at the base of the convective envelope is hot enough (
K) and 12C can be converted into 13C and 14N
through the CN cycle (Sackmann & Boothroyd 1992; Wood et al. 1983).
The HBB models (Sackmann & Boothroyd 1992; D'Antona & Mazzitelli 1996;
Mazzitelli et al. 1999) also predict the production of the
short-lived 7Li by the chain 3He(
,
)7Be
(e-,v)7Li, through the so-called "7Be transport mechanism''
(Cameron & Fowler 1971). One of the predictions of these models is that
lithium should be detectable, at least for some time, on the stellar surface.
Activation of HBB in massive O-rich AGB stars is supported by studies of AGB stars in
the Magellanic Clouds (hereafter, MCs) (Plez et al. 1993; Smith & Lambert
1989, 1990a; Smith et al. 1995), which show a lack of high-luminosity C stars beyond
.
Instead, the more luminous AGB stars in the MCs are O-rich.
Detection of strong Li overabundances, together with strong s-element enhancement in
these luminous AGB stars in the LMC and in the SMC, is the signature that these stars
are indeed HBB AGB stars that have undergone a series of thermal pulses and dredge-up
episodes in their recent past. The HBB nature of these stars is also confirmed by the
detection of a very small 12C/13C ratio (
3-4), expected only when HBB is
active (Mazzitelli et al. 1999). Unfortunately, current HBB theoretical
models have been tested almost exclusively using the results of studying the more
massive AGB stars in the MCs but have never been applied to Galactic sources, mainly
because of the lack of observations available but also because of inaccurate
information on distances within our Galaxy.
Indeed, only a handful of Li-rich stars have been found in our Galaxy so far (e.g. Abia et al. 1991, 1993; Boffin et al. 1993) and, unlike those detected in the MCs,
they are not very luminous (
). Most of them are
low-mass C-rich AGB stars (Abia & Isern 1996, 1997, 2000) and intermediate-mass S- and
SC-stars (Abia & Wallerstein 1998) and not O-rich M-type stars. A few are less
luminous red giant branch (RGB) stars at the bump (Charbonnel 2005). Some of these AGB
stars are among the most Li-rich stars in our Galaxy, the so-called "super Li-rich''
AGB stars (Abia et al. 1991). Under these conditions, however, HBB is not expected to
be active and the Li production is not well-understood. Note that HBB is expected to be
active only in the most massive (and luminous) AGB stars (from
4 to 7
)
(Mazzitelli et al. 1999), which should not be C-rich, but
O-rich. Actually, the best candidates in our Galaxy are the so-called OH/IR
stars, luminous O-rich AGB stars that are extremely bright in the infrared, showing a
characteristic double-peaked OH maser emission at 1612 MHz. These stars are also known
to be very long period variables (LPVs), sometimes with periods of more than 500 days
and large amplitudes of up to 2 bolometric magnitudes. However, they experience very
strong mass loss rates (up to several times 10-5
yr-1) at the end
of the AGB, and most of them appear heavily obscured by thick circumstellar envelopes,
making optical observations very difficult (e.g. Kastner et al. 1993; Jiménez-Esteban
et al. 2005a, 2006a). Thus, no information exists yet on their lithium abundances
and/or possible s-process element enrichment.
In this paper we present results from a wide observational programme based on high-resolution optical spectroscopy of a carefully selected sample of Galactic AGB stars thought to be massive from their observational properties. The criteria followed to select the sources in the sample are explained in Sect. 2. The high-resolution spectroscopic observations made in the optical and the data reduction process are described in Sect. 3, while the main results are presented in Sect. 4. In this section we also show how spectral synthesis techniques were applied to derive the atmospheric parameters, as well as the lithium and zirconium abundances, of the stars in our sample. A discussion of these results in the context of HBB models, nucleosynthesis models, and stellar evolution can be found in Sect. 5. Finally, the conclusions derived from this work are given in Sect. 6.
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Figure 1: Galactic distribution of the sources in the sample. |
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Figure 2: IRAS two-colour diagram [12] - [25] vs. [25] - [60], where the positions of all the sources in the sample are represented (black dots). The continuous line is the "O-rich AGB sequence'' (see text) which defines the sequence of colours expected for O-rich AGB stars surrounded by envelopes with increasing thickness and/or mass loss rates (Bedijn 1987). |
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Stars were included in the sample if they satisfied at least one of the above
conditions (ideally as many of them as possible), which guarantees that they
are actually relatively massive stars. In Table 1 we list the 102 stars
selected for analysis, together with their IRAS names and other being names from
the literature, the Galactic coordinates, the variability type taken from the
Combined General Catalogue of Variable Stars (GCVS, Kholopov 1998), the
spectral type taken from Kwok et al. (1997) and references therein,
and the IRAS colour indices [12] - [25] and [25] - [60]. Similar information for a
few well-known Galactic M-supergiants and S-, SC-, and C-type AGB stars that
were also observed for comparison purposes is presented in Table 2. In
addition, nine stars classified as Mira-type stars in SIMBAD
and included in the initial sample were found to
exhibit observational properties corresponding to C-rich stars. These are
listed in Table 3.
The Galactic distribution of the stars in our sample is shown in Fig. 1. The distribution observed clearly suggests that most of them must belong to the disc population. A few sources with bright optical counterparts located at high Galactic latitudes may represent a subset of nearby stars.
In Fig. 2 we show the position of these stars in the IRAS two-colour diagram [12] -[25] vs. [25] - [60]. They fall along the observational "O-rich AGB sequence'' (García-Lario 1992), as was expected and with very few exceptions. This sequence agrees very well with the model predictions by Bedijn (1987) and is interpreted as a sequence of increasing thickness of the circumstellar envelopes and/or mass loss rates in O-rich AGB stars.
We used a TEK 1124
1124 CCD detector during the first run at the
4.2 m WHT with UES in August 1996 and a SITe1
CCD during the
second and third runs in June 1997 and August 1997, respectively. Since we were
mainly interested in the spectral range between 6000 Å and 8200 Å, we used
the 31.6 line/mm grating in order to provide full coverage of this spectral
region in a single spectrum. With a central wavelength around 6700 Å our
spectra extended over
4000 Å during the first run and
6000 Å during the second and third runs. The spectra were taken with a separation
between orders of around 21 pixels (or 7.5''), which is large enough to allow sky
subtraction, taking into account that the targets in our sample are
non-extended. The resolving power was around 50 000, equivalent to a spectral
resolution of 0.13 Å around the Li I line at 6708 Å. The selected set-up
covered the spectral regions from 5300 Å to 9400 Å in about 47 orders with
small gaps in the redder orders for the first run while the 4700-10 300 Å region was covered in about 61 orders without any gaps for the second and
third runs.
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Figure 3:
High-resolution optical spectra of sample stars displaying
the lack of the ZrO absorption bands at 6474 and 6495 Å compared with two
Galactic S-stars (WY Cas and R And), where these bands are prominent.
WX Ser (IRAS 15255+1944) and V697 Her
(IRAS 16260+3454) are Li-detected while S CrB (IRAS 15193+3132) and IRAS
16037+4218 are Li undetected. The absorption band
at ![]() |
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CASPEC spectra, taken at the ESO 3.6 m telescope, covered the wavelength range
from 6000 Å to 8200 Å. The red cross-disperser (158 line/mm) was used,
which gave a resolving power of 40 000 (equivalent to a spectral
resolution of
0.17 Å around the Li I line at 6708 Å) and an adequate
interorder separation, using the TEK
CCD. With
this set-up, the selected spectral region is completely covered in 27 orders
with small gaps only between the redder ones.
The observational strategy was similar in all runs. Since the stars of the sample are
known to be strongly variable, typically 3-4 mag in the R-band and sometimes
more than 8 mag in the V-band, we determined the exposure time according to
the brightness of the source at the telescope. The typical exposure times ranged from 10 to 30 min. Two 30-min spectra were taken only for a few very faint stars (and
later co-added in order to increase the S/N ratio of the final spectrum). The goal was
to achieve an S/N ratio of 50-150 in the region around 6708 Å (the lithium line) in
order to resolve this narrow absorption line, usually veiled at these wavelengths by
the characteristic molecular bands which dominate the optical spectrum of these
extremely cool stars. Note that, because of the very red colours of the sources
observed, the S/N ratios achieved for a given star can strongly vary from the blue to
the red orders (e.g. 10-20 at 6000 Å while >100 at
8000 Å).
At the telescope, all the 102 O-rich stars listed in Table 1 were tried, but useful spectra were obtained only for 57 (56%). The remaining 45 (44%) sources were either too red to obtain any useful information on the strength of the Li I line at 6708 Å or the optical counterpart was simply not found, in the most extreme cases. Some targets with known, relatively bright optical counterparts in the DSS plates were too faint to be detected on the date of the observations. In contrast, other targets with no optical counterpart in the DSS plates appeared as considerably bright stars. A few Galactic M-supergiants together with some C-, SC-, and S-type AGB stars were also observed for comparison (see list in Table 2). Some of the stars in this latter group are among the most Li-rich stars in our Galaxy. In addition, as we have already mentioned in Sect. 2, a few stars in the initial sample turned out to be carbon stars when observed, and not O-rich as initially suspected (see list in Table 3). The total number of objects observed was 120.
For wavelength calibration, several Th-Ar lamp exposures were taken every night during the UES runs. In the case of the CASPEC observations, these calibration lamp exposures were taken before or after any science exposure at the position of the target in order to keep control of possible dispersion changes with the telescope position. Note that with UES the system is more stable because the instrumentation is located at the Nasmyth focus of the 4.2 m WHT telescope. Finally, the corresponding bias and flatfield images were also taken at the beginning (or at the end) of the night.
The two-dimensional frames containing the echelle spectra were reduced to
single-order one-dimensional spectra using the standard ECHELLE software
package as implemented in IRAF.
Basically, the data reduction process consists of bias-level and
scattered-light subtraction, the search and trace of apertures using a
reference bright star, the construction of a normalized flatfield image to
remove pixel-to-pixel sensitivity fluctuations as well as the extraction of the
1D spectra from the 2D frames. For the wavelength calibration we selected
non-saturated Th-Ar emission lines and third or fourth order polynomials for
the fitting. The calibration accuracy reached was always better than 20 m
.
Finally, we identified the terrestrial features (telluric absorption
lines) comparing the target spectra with the spectrum of a hot, rapidly
rotating star observed on the same night. It should be noted, however, that
the majority of the spectral ranges used in the abundance analysis
presented in this paper are not significantly affected by these features.
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Figure 4: High-resolution optical spectra of sample stars displaying increasing strength of the Li I line at 6708 Å. Note the P cygni-type profile of the Li I line in IRAS 18057-2616. The jumps at 6681 and 6714 Å correspond to bandheads of the TiO molecule. |
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In general, all stars show extremely red spectra with the flux level falling
dramatically at wavelengths shorter than 6000 Å. In addition, the spectra are
severely dominated by strong molecular bands mainly due to titanium oxide
(TiO), as a consequence of the very low temperature and the O-rich nature of
these stars. The bandheads of TiO at
6651, 6681, 6714,
7055, and 7125 Å are clearly present in all spectra. Interestingly, the
bandheads of ZrO at
6378, 6412, 6474, 6495, 6505, and 6541 Å seem to be
absent. These ZrO bandheads (together with those corresponding to other
s-element oxides such as LaO or YO) are very strong in Galactic S-stars (see
Fig. 3).
The TiO veiling effect is so intense that it is very difficult to identify
individual atomic lines in the spectra of these stars with the exception of the
Li I line at 6708 Å, the Ca I lines at 6122 and 6573 Å, the K I line at
7699 Å, the Rb I line at 7800 Å, and a few strong Fe I lines. The K I and Rb I resonance lines sometimes show complex profiles, emission over absorption,
and blue-shifted components which may have their origin in the expanding
circumstellar shell. Actually, the difference between the mean radial
velocities of the stellar and circumstellar line components is found to be of
the order of the expansion velocity derived from the OH maser emission. In a
few cases the Li I line can also show a similar behaviour (see Fig. 4).
Finally, some stars display H emission, which in this type of stars
is generally interpreted as the consequence of the propagation of shock waves
through the outer layers of the stellar atmosphere.
We detected the presence of the Li I resonance line at 6708 Å in 25% of the sources in the sample (25 stars) with a wide variety of strengths, while in 31% of them (32 stars) we did not find any signature of this line. The remaining 44% (45 stars) were either too red or the optical counterpart was simply not found at the moment of the observations. Sample spectra of stars showing increasing strength of the Li I line at 6708 Å are presented in Fig. 4.
The observed sample has been divided into two different groups on the basis of the detection or non-detection of the lithium line in Tables 5 and 6, respectively. A third group is formed by the stars that were too red (or which did not show any optical counterpart at the telescope), and these are listed in Table 7. The IRAS name, the run in which the stellar field was observed, the OH expansion velocity, and the pulsational period (if available) are given for every source in these tables. The same information is also presented for the comparison stars in Table 8 and for the sample of peculiar C-rich AGB stars in Table 9 (in this latter case the expansion velocities are derived from CO data).
The detection of strong circumstellar Li I and Rb I lines in S Per (IRAS 02192+5821), sometimes classified in the literature as an M-type Galactic
supergiant, is remarkable. M-type supergiant stars are not expected to show a
strong Li I line (e.g. Luck & Lambert 1982) or any s-process element
enhancement (e.g. Smith et al. 1995). According to our data, S Per looks like a
possible genuine O-rich AGB star. In addition, we find a cool effective
temperature of 3000 K which is quite different from Levesque et al. (2005)
and further investigation is needed. Similarly, PZ Cas (IRAS 23416+6130),
another source usually classified as an M-type supergiant (e.g. Levesque et al. 2005), is for the first time clearly identified by us as a possible S-,SC-type
AGB star. PZ Cas seems to be slightly enriched in Zr and displays some strong Ba atomic lines but it is not enriched in Li, as expected if PZ Cas is a low-mass
S-,SC-type AGB star. Consistent with our possible AGB identification for the
latter two sources, they show variability amplitudes of more than 3 mag
in the V-band and are the only two stars with periods well beyond 500 days in
Table 8. A more detailed analysis of these two stars is beyond the scope of this
paper but it would be needed to reach a definitive conclusion. Also interesting
is the new detection of a strong lithium line in the peculiar mixed chemistry
(C-rich and O-rich) star IRAS 09425-6040, which is studied in a separate
paper (García-Hernández et al. 2006a). Lithium was also detected in the
C-rich AGB stars IRAS 04130+3918, IRAS 20072+3116 (V1969 Cyg), and IRAS 23320+4316 (LP And), for which a detailed analysis will be presented
elsewhere. Finally, the S-star IRAS 10436-3459 (Z Ant), observed for the
first time with high-resolution spectroscopy, shows, as expected, intense
molecular bands of s-element oxides (like ZrO and LaO) but no lithium. This
star will also be analysed in detail in a separate publication.
Hydrostatic model atmospheres based on the MARCS code (Plez et al. 1992; Gustafsson et al. 2003) reproduce well the observed spectra of C-rich AGB stars (e.g. Loidl et al. 2001) and O-rich AGB stars (e.g. Alvarez et al. 2000) in the optical domain and were thus adopted for the analysis. Indeed, the optical spectra of massive O-rich AGB stars in the MCs have succesfully been modelled using this code (e.g. Plez et al. 1993). A comparison of these models with high-resolution observational data in the optical for massive O-rich AGB stars in our Galaxy is unfortunately still lacking.
Our analysis combines state-of-the-art line-blanketed model atmospheres and
synthetic spectroscopy with extensive line lists. For this we need to estimate
first the range of values of the stellar parameters that can reasonably be
adopted for the stars in our sample: effective temperature
,
surface gravity log g, mass M, metallicity
Z = [Fe/H], microturbulent velocity
,
and C/O ratio.
Spherically symmetric, LTE, hydrostatic model atmospheres were calculated using
the MARCS code for cool stars (Gustafsson et al. 2003). The models are
designed with the notation (
,
log g, M, Z,
,
C/O), where
is the effective temperature of the star in K, g is
the surface gravity in cm s-2, M is the stellar mass in
,
Z is the metallicity,
is the microturbulent
velocity in km s-1, and C/O is the ratio between the C and O abundances.
Synthetic spectra were generated with the TURBOSPECTRUM package (Alvarez &
Plez 1998) which shares much of its input data and routines with MARCS. Solar
abundances of Grevesse & Sauval (1998) were always adopted, except for the
iron abundance, which was taken as log
(Fe) = 7.50.
Most opacity sources for cool M-type stars were included and taken from the
literature. In addition, updated line lists were produced by us for TiO (Plez
1998), ZrO (Plez et al. 2003), and VO (Alvarez & Plez 1998).
For atomic lines, the primary source of information was the VALD-2 database
(Kupka et al. 1999). The NIST
atomic spectra database was also consulted for comparison. Where possible,
gf values were checked by fitting the solar spectrum. VALD-2
gf values taken from Kurucz (1993, 1994) were often erroneous and
did not fit the atomic lines of the solar spectrum. The identification of
features was made using the solar atlases by Moore et al.
(1966) and Wallace et al. (1993, 1998) and the observed solar
spectrum of Neckel (1999). The TURBOSPECTRUM program was run using the
abundances from Grevesse & Sauval (1998) together with the solar model
atmosphere by Holweger & Müller (1974) with parameters
K, log g = 4.44, and a variable microturbulence
as a function of the optical depth.
The whole machinery was tested on the high-resolution optical spectrum
(
)
of the K2 IIIp giant Arcturus (
Boo) from Hinkle et al. (2000). A MARCS model atmosphere was used with the fundamental parameters
determined by Decin et al. (2003a):
K, log g = 1.50,
0.5,
,
km s-1, C/N/O = 7.96/7.55/8.67, and 12C/13C = 7.
These parameters are in excellent agreement with other determinations reported
in the literature. Some Kurucz gf values of metallic lines of Fe I,
Ni I, Co I, Cr I, Ti I, Zr I, and Nd II, which were too weak in the solar
spectrum, were adjusted so as to yield a good fit to the Arcturus spectrum in
several regions (especially in the 7400-7600 Å and 8100-8150 Å regions);
otherwise we used the gf values from the VALD-2 database. This exercise
was very useful in confirming the lack of s-process atomic lines in the
7400-7600 and 8100-8150 Å spectral windows in the Galactic O-rich AGB stars in our sample, as we shall see later.
For the abundance analysis we concentrated our attention on the following spectral regions.
We also synthesized the spectral regions around the resonance lines of Rb I at 7800 Å and K I at 7699 Å, with the intention of using the elemental
abundances derived from these lines as neutron density and metallicity
indicators, respectively. In addition, these spectral regions were also used
to check that the model parameters adopted to reproduce the lithium region
also provided a reasonably good fit at other wavelengths. Two intervals
covering
60 Å in the regions 7775-7835 Å and 7670-7730 Å were
selected for this purpose. Unfortunately, the frequent detection of
circumstellar components to these lines prevented the accurate determination
of the K I elemental abundances, and they will not be discussed here. To a
lesser extent the problem also affects the Rb I line. A detailed analysis of
the Rb abundances will be presented in a forthcoming paper (García-Hernández et al. 2006b).
In order to analyze how the variations in stellar parameters influence the
output synthetic spectra, we constructed a grid of MARCS model atmospheres and
generated the associated synthetic spectra with the following specifications:
i) the stellar mass was in all cases taken to be 2 ;
ii)
values ranging from 2500 to 3800 K in steps of 100 K; iii) C/O ratio values of 0.5, 0.7, 0.8, and 0.9; iv) log g between
-0.5 and 1.6 dex in steps of 0.3 dex; v) microturbulent velocity
between 1 and 6 km s-1 in steps of 0.5 km s-1; vi) metallicity Z between 0.0 and -0.3 dex; vii) log
(Zr)
from 1.6 to 3.6 dex in steps of 0.25 dex; viii) CNO abundances shifted
1.0 dex (in steps of 0.5 dex) from the solar values of Grevese & Sauval
(1998); and ix) 12C/13C ratios of 10 (as expected for HBB stars) and 90 (the solar value). Finally, the synthetic spectra were convolved
with a Gaussian profile (with a certain FWHM typically between 200 and 600 mÅ) to account for macroturbulence as well as instrumental profile
effects.
This, however, results in an extremely large grid of synthetic spectra containing thousands of possible combinations of the above parameters! In order to reduce the number of spectra to be considered in our analysis, further constraints were imposed by studying the sensitivity of our spectral synthesis to changes in these parameters.
The synthetic spectra are, in contrast, particularly sensitive to
,
which determines the overall strength of the molecular
absorption over the continuum (mainly, from TiO molecules but also from VO, and
from ZrO if the Zr elemental abundance is increased above a certain limit).
Thus, we decided to keep the whole range of values initially considered in our
grid, i.e.
from 2500 to 3800 K in steps of 100 K.
Actually, the depth of the molecular bands increases considerably with
decreasing temperature. The effective temperature has also a large impact on
the strength of the Li I line at 6708 Å in the synthetic spectrum. A
decrease of
has the effect of increasing the TiO veiling,
as well as the strength of the Li I line, as a consequence of the changes in
the Li I/Li II population equilibrium.
The C/O ratio basically determines the prevalence of O-rich molecules
C/O < 1) against C-rich molecules like CN, C2, etc. In the case
of our O-rich AGB stars, the precise value of the C/O ratio adopted,
always <1, affects mainly the strength of the TiO veiling and of the Li I
line in the synthetic spectrum. An increase of the C/O ratio will
decrease the TiO veiling and increase the Li I line strength relative to the
adjacent continuum in a very similar way as a decrease in
.
This is a good example of how a set of parameters providing a good match to the
observations for a given spectral region is not necessarily unique. A good fit
may be obtained for another combination of parameters that can correspond to
quite different Li abundances. Fortunately, one of these two sets of
parameters usually does not provide acceptable fits for other spectral
regions. In particular, the presence of detectable atomic lines in the
synthetic spectrum is very sensitive to variations in the C/O ratio:
an increase in the C/O ratio will lower the TiO veiling, making the
detection of atomic lines easier. According to the models, the immediate effect should
be the detection of much stronger atomic lines of K I, Fe I, Zr I, Nd II, etc.
Since we do not see any such effect in our spectra, we conclude that C/O must always be
0.75 in our stars. Abundances derived from models with
C/O between 0.15 and 0.75 show actually little differences (Plez et al. 1993). Given that all stars in our sample are clearly O-rich
and taking into account the above considerations, we decided to use
in all cases, as a constant value. This selection is in
agreement with other more detailed determinations made in massive O-rich AGB
stars in the MCs (e.g. Plez et al. 1993; Smith et al. 1995) and in a
few low-mass M-type stars studied in our Galaxy (Smith & Lambert 1985,
1990b)
.
The surface gravity also affects the appearance of the output synthetic spectra
but its effect is small compared to that of the effective temperature. The
influence of a change in the value adopted for the surface gravity within the
range of values (log g between -0.5 and 0.5 dex) under consideration
is not as severe as a temperature change. The molecular absorption becomes
slightly weaker at higher surface gravities. For example, an increase in the
surface gravity of 0.5 dex has approximately the same impact as an increase of
100 K in effective temperature. However, for a fixed temperature we found that
the lowest gravities in general fit better both the TiO band strength and the
pseudo-continuum around the 6708 Å Li I line. As the appearance of the spectra
is not so dependent on the surface gravity, and considering that its value must
be low in these mass-losing stars, a constant surface gravity of log g = -0.5 was adopted for all the stars in the sample. This selection seems to
be also appropriate if it is compared with the values adopted by Plez et al. (1993) of
for O-rich AGB stars
in the MCs with effective temperatures between 3300 and 3650 K.
The microturbulent velocity is usually derived by demanding no correlation
between the abundance of iron derived from individual lines and their reduced
equivalent widths. Previous studies of O-rich AGB stars found microturbulent
velocities between 2 and 4 km s-1 (Plez et al. 1993; Smith &
Lambert 1985, 1989; Smith et al. 1995; Vanture & Wallerstein 2002).
Unfortunately, we cannot estimate the microturbulent velocity in the stars of
our sample due to the lack of useful Fe I atomic lines (the small number of
detectable Fe I lines does not cover a wide range of equivalent widths). Thus,
we will assume in the following a microturbulent velocity km s-1,
which is a typical value generally adopted for AGB stars (e.g. Aringer et al. 2002). Higher values of
would strenghten all
atomic lines in the synthetic spectra and we do not see this effect in our
data.
The metallicity of the Galactic O-rich AGB stars in our sample is assumed to
be solar, as expected for stars belonging to the disc population of our Galaxy.
A lower metallicity is unlikely since all the stars are expected to be at least
of intermediate mass (
), if not more massive. This
assumption is also in good agreement with the typical metallicities derived
for other AGB stars in our Galaxy (e.g. Abia & Wallerstein 1998; Vanture &
Wallerstein 2002). The Ca I lines at
6122 Å and 6573 Å generally used
for this determination in other studies are unfortunately not sensitive enough
to metallicity nor to surface gravity at the very low temperature of
our sample stars (see Fig. 7 of Cenarro et al. 2002). This was checked
by running different test models and spectral syntheses in the 6100-6160 Å and 6535-6585 Å regions where these lines fall. A decrease in the
metallicity implies less availability of metals like Ti, and thus, lowers the
TiO veiling, making the detection of atomic lines easier. Actually, this may be
the main reason why they are easily detected in O-rich AGB stars at the
metallicity of the MCs (e.g.
0.5 in the SMC, Plez et al. 1993), while we do not detect them in our spectra.
After imposing these further constraints, the grid of MARCS model spectra to
consider is composed of "only'' a few hundred spectra, with effective
temperatures in the range
K in steps of 100 K,
log
(Zr) between 1.6 and 3.6 dex in steps of 0.25 dex,
and a variable value of the FWHM in the range 200-600 mÅ in steps
of 50 mÅ, keeping all the other stellar parameters fixed: log g = -0.5,
,
solar metallicity Z= [Fe/H]=0.0,
km s-1, C/O = 0.5, solar CNO abundances, and solar
12C/13C ratios.
In order to find the synthetic spectrum which better fits the observed spectrum
of a given star a modified version of the standard 2 test was
used. When fitting observed data
to model data
,
the quality of the fit can be quantified by the
2 test. The best fit corresponds to that leading to the minimum
value of
2. This test is used here in a modified way and is
mathematically expressed as:
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(1) |
The observed spectra, once fully reduced, were shifted to rest-wavelength using
the mean radial velocity shift derived from the Li I and Ca I lines. In
addition, they were re-binned to the same resolution as the synthetic ones
(0.06 Å/pix) and normalized in order to make the comparison easier. The
observed spectra were then compared to the synthetic ones in intervals of 60 Å which allowed the analysis of their overall characteristics as
well as of the relative strength of the TiO molecular bands.
The fitting procedure makes a special emphasis on the goodness of the fit in
the lithium region (6670-6730 Å). In this spectral range the goal was to
fit the relative strength of the TiO bandheads at 6681 and 6714 Å,
which are very sensitive to the effective temperature, together with the
pseudo-continuum around the Li I line. We first determined by
2 minimization which of the spectra from our grid of models provided the best
fit, mainly fixing
.
The best fits resulting from the
2 minimization were also judged by eye in order to test the method.
The lithium content was then estimated by changing the lithium abundance. This
procedure was repeated on each star of the sample for which an acceptable S/N ratio was achieved around the Li I 6708 Å line. Unfortunately, we could not
analyse a few AGB stars with a low signal-to-noise ratio in their spectra at 6708 Å (see Table 7). IRAS 18025-2113, IRAS 03507+1115, and IRAS 18304-0728 showed unusual spectra, very different from the rest of the stars
in our sample and a good fit from our grid of MARCS model spectra was not
possible. The spectra of these stars show unusual TiO band strengths which
could only be reproduced with a low
but the rest of the
spectrum (e.g. the local continuum level) seems to be hotter. In addition, the
atomic lines are broader (IRAS 18025-2113 and IRAS 18025-2113) or narrower
(IRAS 03507+1115) than the FWHM needed to describe the TiO bandheads
and lines. Their complicated spectra look like a combination of two
temperatures and FWHMs suggesting that they could probably be
double-lined spectroscopic binaries or strongly affected by shock waves propagating
in their atmosphere.
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Figure 5:
Best model fit and observed spectrum around the Li I line (6708 Å)
for the star IRAS 11081-4203. The Li abundance derived from this spectrum was
log
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Figure 6:
Best model fit and observed spectrum in the region 6670-6730 Å for
the star IRAS 11081-4203. The
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Figure 7:
Best model fit and observed spectrum in the region 6460-6499 Å for the star IRAS 11081-4203. The Zr abundance derived from this
spectrum was log
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Figure 8:
Best model fit and observed spectrum in the region 7670-7730 Å (around the K I line at 7699 Å) for the star IRAS 11081-4203. The
synthetic spectrum obtained for solar K abundance (log
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Figure 9:
Best model fit and observed spectrum in the region
7775-7835Å (around the Rb I line at 7800 Å) for the star IRAS
11081-4203. The synthetic spectra obtained for solar Rb abundance
(log
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The best model fit in the Li I spectral region (6670-6730 Å) is also usually found to fit the wavelength regions around the ZrO bandhead and the K I and Rb I resonance lines reasonably well. In general, the overall shape of the spectrum (including the TiO bandheads) is very well reproduced. As an example, the fits made in different spectral regions for the star IRAS 11081-4203 are presented in Figs. 5 to 9. The effective temperatures of the best-fitting model spectra together with the lithium abundances (or upper limits) derived following the above procedure are listed in Table 10, where we have separated the Li detected stars from the Li undetected ones.
For the measurement of the zirconium abundance we concentrated our attention on the synthesis of the 6455-6499 Å spectral region. In this case, the goal was to obtain the best fit around the ZrO molecular bands at 6474 Å and 6495 Å. The non-detection of these features in any of the stars under analysis imposes severe upper limits to the zirconium abundance. To be sure that there was no other way to interpret our data, we also checked the presence of atomic zirconium in the 7400-7600 and 8100-8150 Å spectral regions, where strong Zr I atomic lines should have been detected in that case. Unfortunately, the local continuum in these latter wavelength regions is not well reproduced in the O-rich AGB stars in our sample, in contrast with the perfect agreement obtained in some of the comparison stars observed, such as the S-star IRAS 10436-3454, for which a perfect fit is obtained both for the ZrO bandheads at 6455-6499 Å (see Fig. 12) as well as for the Zr I atomic lines in the 7400-7600 and 8100-8150 Å spectral regions (not shown). Sample spectra of the region around 7565 Å are shown in Fig. 10, where the position of some atomic lines of Zr I, Nd II, and Fe I are indicated. As we can see, the atomic lines corresponding to s-process elements (e.g. Zr, Nd, La, etc.) were not detected in any of the sample stars observed. In contrast, these s-element atomic lines appear very strong in the Galactic S-star IRAS 10436-3454. Note, however, that the stars in our sample are the coolest O-rich AGB stars ever studied (with effective temperatures between 2700 K and 3300 K; see Table 10). We suspect that the strong discrepancies observed between model and observations may be due to the non-inclusion of other O-rich molecules, such as H2O, in our line lists. The effect of H2O may be ignored in S-stars, where the C/O ratio is close to unity, but it might become dominant in our O-rich stars at these very low temperatures, preferentially at the longer wavelengths (e.g. Allard et al. 2000; Decin et al. 2003a,b). In addition, lanthanum oxide also has a strong absorption band around 7500 Å and it could have been a good candidate for detection in the spectral range here considered. As for ZrO, the non-detection of the strong LaO molecular bandhead, usually very strong in S-stars at these wavelengths, is another indication of the lack of s-process elements in our sample stars.
Table 10:
Spectroscopic
and Li abundances1 derived for the
subgroup of Li detected (left) and Li non-detected stars (right).
The derived Li abundances are displayed in Table 10, where we can see that
almost all Li detected stars show enhanced Li abundances log
(Li
,
i.e. larger than solar, but smaller than those
found in the so-called "super Li-rich'' AGB stars (with log
(Li) > 3-4; e.g. Abia et al. 1991) in our Galaxy. A
very similar range of Li overabundances was found in the massive O-rich AGB
stars studied in the MCs (Plez et al. 1993; Smith & Lambert 1989,
1990a; Smith et al. 1995). The errors in the derived Li abundances mainly
reflect the sensitivity of our models to changes in the adopted stellar
parameters. In particular, they strongly depend on the uncertainties in the
determination of the effective temperature (
100-200 K; larger for the
coolest stars) and metallicity (
0.3), while they are less sensitive to
the microturbulent velocity (
1 km s-1), surface gravity (
0.5),
and FWHM (
50 mÅ) uncertainties. The effect of the uncertainty
in the location of the pseudo-continuum around the Li I line on the measured
abundances is negligible compared with any other change in the adopted stellar
parameters. If we consider all these uncertainties as independent sources of
error, the resulting Li abundances given in Table 10 are estimated to be
affected by errors of the order of 0.4-0.6 dex. As an example, the changes in
the derived Li abundance induced by slight variations of each of the
atmospheric parameters used in our modelling for IRAS 15255+1944 are
shown in Table 11.
Note, however, that the estimated errors do not reflect possible non-LTE effects, dynamics of the stellar atmosphere, or errors in the model atmospheres or in the molecular/atomic linelists themselves. In particular, over-ionization and over-excitation of Li is predicted to occur under non-LTE conditions both in C-rich and O-rich AGB stars (Kiselman & Plez 1995; Pavlenko 1996) although the effect is expected to be more pronounced in metal-deficient stars (i.e. MC AGBs). Use of LTE in Li-rich AGB stars is likely to result in underestimation of the Li abundance (Kiselman & Plez 1995; Abia et al. 1999).
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Figure 10:
High-resolution optical spectra around the 7565 Å region of two
stars in our sample where we show the lack of s-process element
atomic lines (e.g. from Zr and Nd) in comparison with the Galactic S-star IRAS
10436-3459. IRAS 05027-2158 and IRAS 19361-1658 show lithium lines with
very different strength. The positions of some atomic lines of Zr, Nd, and Fe
are indicated. It seems that only the Fe I line at ![]() |
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Concerning s-process elements, the non-detection of the ZrO molecular bands at 6474 Å and 6495 Å is found to be consistent with upper limits for the Zr abundance around [Zr/Fe
dex with respect to the solar value of
log
(Zr) = 2.6 dex. The synthetic spectra predict
detectable ZrO bandheads at 6474 Å and 6495 Å for Zr abundances above
0.00-0.25 dex for
K and above 0.25-0.50 dex
for
K. Actually, the fits are also reasonably
good even if we do not include any ZrO in the synthesis. We found the same
results both in the Li non-detected stars and in the Li detected ones. This
result is in strong contrast with the higher Zr abundances ([Zr/Fe] > 0.5)
found in Galactic MS-, S-stars, and in massive O-rich AGB stars in the MCs. The
effect of varying the Zr abundance for the Li detected star IRAS 11081-4203
is displayed in Fig. 11, as an example. The best model spectrum fit for the
Galactic S-star IRAS 10436-3454 is also shown in Fig. 12 for comparison.
As we can see, strong ZrO bandheads are detected in IRAS 10436-3454, while
these are completely absent from the spectrum of IRAS 11081-4203. According
to our models, a modest Zr enhancement (with respect to the solar value) would
be enough for the ZrO molecular bands to show up in our sample stars despite
the strong veiling produced by the TiO molecule (see Fig. 11). An even lower
Zr abundance is enough to produce strong ZrO bands in S-stars, where the
C/O ratio is
1 because of the weaker TiO veiling, as we can see
in the spectrum of IRAS 10436-3454 displayed in Fig. 12. This also favours
the visibility of the Zr I atomic lines in the 7400-7600 and 8100-8150 Å spectral regions in these stars, as has already been mentioned.
In constrast to the studies made in the past of AGB stars in the MCs, for which a
common distance can be assumed leading to relatively well-determined absolute
luminosities, the estimation of absolute luminosities for the stars in our Galactic
sample is very difficult, mainly because of the large uncertainties involved in the
determination of distances within our Galaxy. As a first
approach, we can try to estimate luminosities using the period-luminosity
relationship (for those sources with a well determined period).
If the P-L
relationships for LMC LPVs found by Hughes & Wood (1990) are
applied to our sample of OH/IR stars, we obtain
between
-5 and -6.2
for
those stars with periods of 400-600 days while an
of
-4.6 is
derived for periods around
300 days.
However, if we take those stars with the longer periods
in the sample (
1000 days), we derive
7.9, which
seems unrealistically high.
This suggests that the period-luminosity relationship may
not hold for the most extreme OH/IR stars (Wood et al. 1998). On the other
hand, there is recent observational evidence for the existence of low metallicity
(e.g. in the LMC) HBB AGB stars with luminosities brighter than the predictions of the
core-mass luminosity relation that have been atributed to an excess flux from HBB
(Whitelock et al. 2003). A similar effect could explain the detection of AGB stars
with such high
also in our Galaxy.
Determining the progenitor masses of our Galactic O-rich AGB stars is not a simple task either. Several observational parameters, however, such as the OH maser expansion velocity or the variability period, have been proposed in the past as useful distance-independent mass indicators for this class of stars, and we will make use of them in the following.
The assumption is based on the fact that the group of OH/IR stars with longer periods and larger expansion velocities shows a Galactic distribution that corresponds to a more massive population (Baud et al. 1981; Baud & Habing 1983; Chen et al. 2001; Jiménez-Esteban et al. 2005b) suggesting that the periods of variability and the expansion velocities of the circumstellar envelopes of AGB stars must be closely correlated with the progenitor mass.
Table 11: Sensitivity of the derived Li abundances (in dex) to slight changes in the model atmosphere parameters for IRAS 15255+1944.
Actually, stars in our sample follow the expected trend, as is illustrated in Fig. 13. Although the limited number of stars considered prevents us from deriving any statistical conclusion based on our data alone, it seems clear that the Galactic latitude distribution of the stars in our sample becomes narrower as a function ofFigure 14 (left panel) shows the distribution of the sources in our sample on the IRAS two-colour diagram [12] - [25] vs. [25] - [60] for which OH expansion velocities are available. As we can see in this diagram, on average, the sources with intermediate OH expansion velocities, between 6 and 12 km s-1, show redder [12] - [25] and [25] - [60] colours than those with low OH expansion velocities (i.e. below 6 km s-1), which implies thicker circumstellar envelopes or larger mass loss rates, but they are bluer than the subsample of stars with OH expansion velocities beyond 12 km s-1. Redder IRAS colours can be interpreted as the consequence of a more advanced stage on the AGB if we assume that mass loss rate increases along this evolutionary phase. Alternatively, they may also be interpreted as corresponding to a more massive population of AGB stars since in principle they should be able to develop thicker circumstellar envelopes even at a relatively early stage in their AGB evolution. A similar behaviour is observed when the variability period is analysed (Fig. 14; right panel), although the number of stars with known periods is significantly smaller. In this case we can see that the reddest sources in the diagram belong to the subgroup showing the longer periods, in excess of 700 days.
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Figure 11:
Synthetic and observed spectra in the region 6460-6499 Å for the
star IRAS 11081-4203, one of the Li detected stars in our sample. The
synthetic spectra corresponding to [Zr/Fe] = +0.0, +0.5, and +1.0 dex
(or log
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Figure 12:
Synthetic and observed spectra in the region 6460-6499 Å for the
Galactic S-star IRAS 10436-3454. The synthetic spectra corresponding to
[Zr/Fe] = +0.0, +0.5, and +1.0 dex (or log
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The distribution of variability periods and OH expansion velocities is found to be remarkably different for each of the three subgroups identified in our sample. Histograms showing the results obtained are shown in Figs. 15 and 16, respectively. The same results are presented in a different way in Table 12, where the distribution of sources among the different subtypes is presented as a function of their OH expansion velocity and variability properties.
In Table 12 we can see that most of the AGB stars with the lower OH expansion
velocities (
OH) < 6 km s-1) are among those undetected in
the Li I line (92%) while 44% of stars with OH expansion velocities between 6 and 12 km s-1 (but still observable in the optical) are detected as
Li-strong sources. However, there are also stars with non-detection of lithium (24%) and stars too red or not found (32%) within this group. Finally, a
majority of the stars in the group of heavily obscured sources (too red or not
found) shows OH expansion velocities higher than 12 km s-1 (66%).
Similar statistics are obtained for the variability period. Sources with relatively short periods, i.e. below 400 days, are predominantly non-Li detections (84%), while almost half of the stars showing periods between 400 and 700 days are Li detected (46%). Again, the same observational problems arise when we try to analyse the stars with the longer periods (P > 700 days), since most of them are completely obscured in the optical (88%).
In Fig. 17 we can see that there is actually a clear correlation between the Li abundance and the variability period in our target stars. Remarkably, for periods below 400 days (log P < 2.6), only one star out of 12 is Li detected. In contrast, almost all stars with P > 500 days (log P > 2.7) are Li-rich, the only exception being IRAS 18050-2213 (=VX Sgr), which is a peculiar, very cool semi-regular pulsating star, not showing a proper Mira variability in the last decade, which some authors identify as a massive red supergiant, and not as an AGB star (Kamohara et al. 2005).
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Figure 13:
Galactic latitude distribution of the sources
in our sample as a function of their OH expansion velocities:
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Figure 14: Distribution of the sources included in our sample in the IRAS two-colour diagram [12]-[25] vs. [25]-[60] showing the location of the stars in our sample as a function of the OH expansion velocities ( left panel) and variability periods ( right panel). The continuum line is the "O-rich AGB sequence'' (see text), which indicates the sequence of colours expected for O-rich AGB stars surrounded by envelopes with increasing thickness and/or mass loss rates (Bedijn 1987). |
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At present, it is generally accepted that lithium production in massive O-rich
AGB stars is due to the activation of hot bottom burning
at the bottom of the convective mantle, which prevents the formation of
luminous C-rich AGB stars (see details in Sect. 1). If the detection of
lithium is taken as a signature of HBB and the period and the OH expansion
velocity are accepted as valid mass indicators, our results indicate that there
are no HBB stars among the subgroup of AGB stars in our sample with periods
lower than 400 days and
(OH) < 6 km s-1, while most of the AGB stars with periods beyond 500 days and high expansion velocities are quite
probably developing HBB.
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Figure 15: Distribution of variability periods for the Li non-detected ( left panel), Li detected ( middle panel), and too red/not found ( right panel) sources. |
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Figure 16: Distribution of OH expansion velocities for the Li undetected ( left panel), Li detected ( middle panel), and too red/not found ( right panel) sources. |
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Table 12: Distribution of sources among the different subgroups identified in the sample as a function of their OH expansion velocity and variability properties.
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Figure 17: Observed Li abundance vs. variability period. Upper limits to the Li abundance are shown with black triangles and marked with arrows. Abundance values derived from Li I lines which are resolved in two components (circumstellar and stellar) are shown with red triangles and correspond to the photospheric abundance needed to fit the stellar component. The dashed vertical line corresponds to a variability period of 400 days (see text). |
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Figure 18: Observed Li abundance vs. OH expansion velocity. Upper limits to the Li abundance are shown with black triangles and marked with arrows. The abundance values derived from Li I lines which are resolved into two components (circumstellar and stellar) are shown with red triangles and correspond to the photospheric abundance needed to fit the stellar component. The dashed vertical line corresponds to an OH expansion velocity of 6 km s-1 (see text). |
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Modelling of HBB started thirty years ago with envelope models including
non-instantaneous mixing coupled to the nuclear evolution (Sackmann et al. 1974). More recent models (Sackmann & Boothroyd 1992; Mazzitelli et al. 1999, hereafter SB92 and MDV99, respectively) seem to work
well in reproducing the lithium production observed in massive O-rich AGB stars in
the MCs. Although the precise determination of the lithium production depends
on the input physics (stellar mass, metallicity, mass loss rate, overshooting,
etc.), the strongest dependence is with the convection model assumed. SB92 models predict lithium production in intermediate-mass stars assuming the
so-called "mixing length theory'' (MLT) framework for convection. But these
models require a fine tuning of the mixing length parameter .
MDV99
explored lithium production in intermediate-mass stars according to a more
modern treatment of turbulent convection, known as the "Full Spectrum of
Turbulence'' (FST) model (D'Antona & Mazzitelli 1996, and references therein).
Actually, MDV99 models are able to reproduce the strong Li enhancements
observed in massive O-rich LMC AGB stars in a satisfactory way and provide a
theoretical explanation to the existence of a few sources in the LMC with
extreme luminosities of up to
7.3 and -7.6, which are known
to be long period, obscured AGB stars (Ventura et al. 2000).
However, these models have not yet been tested against their Galactic analogues.
For solar metallicity, MDV99 predict lithium production as a consequence of
the activation of HBB for stellar masses 4.5
without core
overshooting and for
including core overshooting. The
amount of lithium produced during the AGB does not depend on the assumed
initial lithium abundance as the star structure loses all memory of the
previous history of lithium at the beginning of the AGB phase. As an example,
according to this model the lithium surface abundance drops to log
(Li
for a 6
star after the second
dredge-up, at the beginning of the AGB phase. If overshooting from below the
convective envelope is allowed, an important decrease in the number of thermal
pulses is predicted by the model and lithium production takes place before the
first thermal pulse. The variation with time of the lithium abundance, total
luminosity, and temperature at the base of the convective envelope (
)
for different masses is shown in Fig. 15 of MDV99. A stronger
Li-overabundance, faster increase of the luminosity at the beginning of the AGB phase, and higher
are predicted when the mass of the star is
increased.
Mass loss rates and metallicities play also an important role according to
MDV99. The minimum mass needed for the ignition of HBB decreases with
decreasing mass loss rates. In addition, a low mass loss rate leads to a longer
run of thermal pulses and to a lithium abundance which remains large for the
entire thermal pulsing AGB (TP-AGB, hereafter) lifetime. In contrast, for high
mass loss rates, the HBB process is stopped only after a few pulses (5
for a 6
star). Models with strong mass loss show oscillations of the
Li abundance by orders of magnitude on short timescales of
104 years
(even within a single thermal pulse!). The temperature at the bottom of the
convective envelope,
,
varies so much that the lithium production is
temporarily stopped while the lithium already produced has already been diluted
by convection. When
increases, the lithium abundance increases
again. The consequence is that, at least for 20% of the time, there is
negligible lithium in the envelope (at a level of approximately log
(Li
). This implies that the distribution of Li abundances derived from the observations can only be analysed in a statistical way.
On the other hand, the activation of HBB takes place at a lower mass limit of
only 3.0-3.7
(depending on the mass loss rate considered) at the
metallicity of the LMC (Ventura et al. 2000). Low metallicity
models predict larger values of
and higher luminosities, and as a
consequence, higher lithium production for a given stellar mass. In addition,
this accelerates Li-production and a maximum value of log
(Li
is reached in a shorter timescale with respect to
the solar metallicity case. The variation with time of the surface lithium
abundance, stellar luminosity, and
,
computed for a 6
model
and different metallicities is shown in Fig. 10 of MDV99. Lower metallicity
models esentially produce lithium faster due to the larger temperature at the
base of the convective envelope, but show high Li abundances for a shorter
time.
In order to interpret these results, several scenarios can be proposed:
The answer to this question must be related to the different metallicity of the stars in the MCs with respect to our Galaxy. Actually, theoretical models predict a higher efficiency of the dredge-up in low metallicity atmospheres (e.g. Herwig 2004) with respect to those with solar metallicity (e.g. Lugaro et al. 2003).
Moreover, there is observational evidence that lower metallicity environments are also less favourable to dust production, as is suggested by the smaller number of heavily obscured AGB stars in the MCs (Trams et al. 1999; Groenewegen et al. 2000), compared to our Galaxy. This is supported by the lower dust-to-gas ratios derived by van Loon (2000) in the few obscured AGB stars in the MCs for which this analysis has been made.
If mass loss is driven by radiation pressure on the dust grains, it might be
less efficient with decreasing metallicity (Willson 2000). In that case, longer
AGB lifetimes would be expected, which could increase the chance of
nuclear-processed material reaching the stellar surface. This would also
explain the larger proportion of luminous C-rich stars (up to
6) observed in the MCs (Plez et al. 1993; Smith &
Lambert 1989, 1990a; Smith et al. 1995) with respect to the Galaxy. The slow
evolution predicted for AGB stars in the MCs as a consequence of the less
efficient mass loss leaves time for more thermal pulses to occur during the
AGB lifetime and, therefore, a more effective dredge-up of s-process elements
to the surface can be expected before the envelope is completely gone at the
end of the AGB. This would explain why even the more massive stars in the MCs
show a strong s-process enrichment in contrast to their Galactic counterparts.
In our Galaxy the only AGB stars showing a similar overabundance in s-process
elements seem to be the result of the evolution of low- to intermediate-mass
stars (
), while no, or very little, s-process
enhancement is observed in Galactic AGB stars with higher main sequence
masses.
In addition, the lower critical mass (which also decreases with decreasing mass
loss rate, see Sect. 5.4) needed to develop HBB (
at
the metallicity of the LMC, compared to the
4
limit in our
Galaxy) favours the simultaneous detection of s-process elements and Li
enrichment in a larger number of AGB stars in the MCs. In our Galaxy, the
Li-rich sample of AGB stars is restricted to the fraction of stars with main
sequence masses
.
In contrast to their MC counterparts,
these stars evolve so rapidly (because of the strong mass loss) that there is
no time for a significant enhancement in s-process elements, as the results
here presented seems to demonstrate.
Obviously, more refined s-process calculations taking into account the effects of HBB and mass loss, using a large grid of stellar evolutionary models with different masses and metallicities are strongly encouraged.
The well known existence of at least two different chemical branches of AGB stars (C-rich and O-rich) and the dependence of this chemical segregation on the main sequence mass and on the environmental conditions in which these stars evolve (such as the metallicity) is a long debated issue in stellar evolution. Moreover, the connection with the chemical composition observed in Galactic post-AGB stars and/or well evolved PNe (which are expected to be the result of the evolution of AGB stars with a wide variety of masses) has never been addressed in detail from an overall perspective.
Stars classified as post-AGB in the literature are mainly low-mass stars,
according to their wide Galactic distribution. They are thought to represent
the fraction of stars which evolve so slowly that they are still visible for
some time as stars with intermediate spectral types before they become PNe.
Many of these stars are C-rich and show strong s-process element enhancements
[s-process/Fe1.5 (van Winckel & Reyniers 2000; Reddy et al. 1999; Klochkova et al. 1999), in agreement with what we would expect if
they were the final product of the evolution of low-mass AGB stars. However,
other post-AGB stars do not show the characteristic signatures of the third
dredge-up. In the non-enriched sources the s-process element abundances can be
as low as [Zr/Fe
(Luck et al. 1990; van Winckel
1997). These objects might be stars with even lower masses (
)
which have evolved off the AGB before the third dredge-up could
occur (Marigo et al. 1999). Some of these sources are
high Galactic latitude, hot B-type post-AGB stars (e.g. Conlon et al. 1993;
Mooney et al. 2001) that show strong carbon deficiencies. In the more extreme
cases they will probably never become PNe.
On the other hand, and as we have shown in this article, the more massive AGB stars may also show a similar lack of s-process elements, but in this case as a consequence of the rapid evolution induced by the strong mass loss. These stars are more difficult to observe in the optical as they are expected to evolve as heavily obscured sources from the TP-AGB to the PN stage but may be the progenitors of most Galactic PNe.
Actually, a similar chemical segregation is also observed in Galactic PNe. Classically, they are classified as type I, II or III as a function of their chemical abundance properties (Peimbert 1978). They are expected to cover a wide range of progenitor masses and thus show different chemical signatures as a consequence of their previous passage along the AGB.
In particular, the higher mass fraction of stars (with
)
identified in this paper as stars developing HBB and showing a strong lithium
enhancement are expected to arrive at the PN stage still as heavily obscured
stars. These stars would evolve into N-rich type I PNe, which are
characterized by their large He abundances (He/H > 0.10) and N/O ratios
(Manchado 2003, 2004; Mampaso 2004, and references therein). They are thought
to represent the fraction of the more massive PNe in the Galaxy, as is
confirmed by their strong concentration at low Galactic latitudes. Their strong N overabundances would be consistent with their identification as PNe resulting
from the evolution of HBB AGB stars.
We have classified the sample into three subgroups on the basis of their OH expansion velocity and variability period. From their relative distribution in the IRAS two-colour diagram [12] - [25] vs. [25] - [60] and Galactic distribution, we conclude that they must represent populations of Galactic O-rich AGB stars with different progenitor masses.
By combining MARCS model atmospheres and synthetic spectroscopy with extensive
line lists we were able to derive the fundamental stellar parameters
(
,
log g, C/O,
,
etc.) and obtain the Li and Zr abundances in those stars for which an optical spectrum had been obtained. Our
chemical abundance analysis shows that half of stars show Li overabundances in
the range between log
(Li
0.5 and +3.0 dex.
This is interpreted as a signature of the activation of HBB, confirming that
they are actually massive AGB stars (
). However, these stars
do not show any zirconium enhancement (taken as representative of the
s-process element enrichment), indicating that they behave differently from the
more massive (and luminous) AGB stars in the MCs.
Assuming that the variability period and the OH expansion velocity can be taken
as distance-independent mass indicators and by comparing our results with the
theoretical predictions of HBB and nucleosynthesis models, we conclude that the
O-rich AGB stars in our sample are all very massive (
),
even those non-detected in lithium; but only the more massive ones (
;
with periods longer than 400 days and
(OH)> 6 km s-1) experience HBB. The lack of lithium enrichment found in some of
the sources thought to belong to the group of more massive AGB stars in the
sample is explained as (i) the consequence of the timescale of the Li production
phase being of the order of or slightly smaller than the interpulse time (
104 years) or, alternatively, (ii) by the 3He exhaustion in the
envelope interrupting the Li production, in agreement with the predictions
made by the models.
The group of stars displaying the most extreme observational properties (with
periods sometimes longer than 700 days and
(OH) > 12 km s-1) are believed to represent the more massive AGB stars in our Galaxy.
Unfortunately, these stars are strongly obscured by their thick circumstellar
envelopes so we could not carry out any type of analysis in the optical range.
This very high-mass population should also experience HBB and show strong Li
enhancement but no s-process element overabundances. We propose that the HBB status of these obscured stars can be determined by measuring their 12C/13C ratios in the near-infrared, which can also be used as an HBB indicator.
Our results suggest that the dramatically different abundance pattern found in AGB stars belonging to the MCs and to our Galaxy can be explained in terms of the different metallicity conditions under which these stars evolved. This constitutes strong observational evidence that the chemical evolution of massive AGB stars is strongly modulated by metallicity.
Globally considered, our results can also be used to establish evolutionary connections between TP-AGB stars, post-AGB stars, and PNe. In particular, we suggest that the HBB AGB stars identified in our sample must be the precursors of N-rich type I PNe.
Acknowledgements
DAGH is grateful to J.I. González-Hernández from the Instituto de Astrofísica de Canarias for the adaptation of hiscode used in this work. AM and PGL acknowledge support from grant AYA 2004-3136 and AYA 2003-9499 from the Spanish Ministerio de Educación y Ciencia.
Table 1: The sample of Galactic O-rich AGB stars.
Table 2: The sample of comparison stars.
Table 3: The sample of peculiar carbon AGB stars.
Table 4: Log of the spectroscopic observations.
Table 5: Galactic O-rich AGB stars where the lithium line was detected.
Table 6: Galactic O-rich AGB stars where the lithium line was not detected.
Table 7: Galactic O-rich AGB stars for which no analysis of the lithium line was possible.
Table 8: Comparison stars.
Table 9: Peculiar C-rich AGB stars.