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Astron. Astrophys. 336, 925-941 (1998) 6. The (
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Fig. 10. The (![]() ![]() ![]() |
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Fig. 11. The (![]() |
It is clear that the majority of the Galactic Center objects do not
follow the LMC () relation for optically visible
LPVs and nearly all lie well below it, with offsets of up to 6
magnitudes. This indicates that most of these objects are surrounded
by cool, dense, dusty shells that obscure the central stars in
K. On the other hand, the majority of the Baade's Window LPVs
lie on the LMC (
) relation as expected (since
the distance given by Glass et al. was derived by fitting to the LMC
relation). However, the longest period Baade's Window LPVs, which are
generally the dusty IRAS sources, lie well below the LMC relation.
Similarly, the LMC OH/IR stars have much longer periods than the
optically visible LPVs of similar K magnitude.
In the () diagram, the optically detected LPVs
in Baade's Window and the LMC once again lie on the same relation,
while most of the Galactic Center LPVs show a lower luminosity,
although the offset is smaller than in the (
)
diagram. Low
values were also found by Jones et
al. (1994) and Blommaert et al. (1997) with much smaller samples of
stars. At least some of the scatter in the (
)
diagram is due to the scatter in the values of
.
This scatter could be reduced by obtaining photometry at wavelengths
longer than 3.6 µm (L).
It is usually assumed that low luminosities or, equivalently, very
long periods for LPVs are a demonstration of the effect of mass loss
on pulsation period. Before the superwind phase of mass loss starts,
LPVs of a given metallicity seem to lie on a unique relation as
demonstrated by the LMC and Baade's Window LPVs. Now, for LPVs in
general, the pulsation period varies as
R/M
, where
1.5-2 and
0.5-0.9
depending on whether the LPVs pulsate in the first overtone or
fundamental mode (Fox & Wood 1982; Wood 1990). Once the stellar
mass has been reduced significantly as a result of a superwind, the
pulsation period increases due to the combined effects of the mass
reduction and the increase in mean radius that results from structural
readjustment (Vassiliadis & Wood 1993). In this situation, the
LPVs evolve almost horizontally from their positions as optical
variables in the (
) diagram to the positions
found here for the Galactic Center OH/IR stars. This sample of stars
is the first to clearly demonstrate this effect over a wide range of
luminosities.
The shorter period stars at the top of the envelope of Galactic
Center variable stars in Fig. 11, in particular the newly
discovered variables, do not have the extremely red intrinsic colours
that might be expected if these stars were surrounded by dense
circumstellar shells produced by superwind mass loss. The pulsation
periods of these stars are probably not greatly affected by mass
reduction as described above. Such stars with
K- 1.3 are shown in Figs. 10 and
11 as filled symbols (we use this colour cutoff as it corresponds to
the red limit of the optically-visible LPVs in Baade's Window). For
300 days, these LPVs have
and
values in excellent agreement with those of
LPVs in the LMC and Baade's Window. The short periods of these stars
and the agreement of their luminosities with those of similar LMC LPVs
suggest that they are old and of metallicity up to solar. For example,
47 Tuc with a metal abundance of
0.2 solar has
three Miras of period
200 days while the globular
clusters NGC5927 and NGC6553, which have near-solar metal abundances
and lie in the inner parts of the Galaxy, appear to have Mira
variables with periods near 300 days (Andrews et al. 1974). This
situation is consistent with solar metallicity being attained in the
Galactic Center region during the earliest star formation epoch.
For 300 days, it is clear that the LPVs
without thick dust shells generally fall below the LMC and Baade's
Window period-luminosity relations. What we are seeing here is a clear
demonstration of the fact that the (
) and
(
) relations for LPVs are metallicity dependent,
as predicted by Wood (1990). The long periods in high metallicity LPVs
result from the giant branch being redder than for solar metallicity
stars. Consequently, the stars will have a larger radius and longer
period at a given luminosity. Note, however, that substantial metal
abundance variations are required to make the metallicity dependence
apparent: Wood(1995) found that the (
) relations
for the LMC (metallicity
0.5 solar) and SMC
(metallicity
0.25 solar) were consistent with
each other.
A rough estimate of the metallicity required to produce the long
periods seen for the less extreme stars in Fig. 11 (solid
symbols) may be obtained as follows. Wood (1990) found that for LPVs
pulsating in the fundamental mode (which we will assume here),
(note that this relation was derived for solar
metallicity and below, so we will be extrapolating beyond the range of
tested parameters here). From Fig. 11, it appears that the bluer
Galactic Center stars with
300 days are
0.4 mag. fainter than bulge or LMC stars of the
same period. Ignoring any possible difference in mass, we then find
from the equation above that we require the Galactic Center
metallicity to be
3.6 times higher than the bulge
metallicity in order to get the observed offset between the bulge
stars and the bluest Galactic Center stars.
Another very important feature of the ()
diagram is that the luminosities of the main group of Galactic Center
OH/IR stars extend considerably brighter than the luminosities of the
LPVs in Baade's Window (we exclude the 4 Baade's Window stars that lie
above the LMC line in Fig. 11 as they are almost certainly
foreground objects). Since final AGB luminosities increase with
initial mass (eg. Vassiliadis & Wood 1993) we interpret this to
mean that there is a significant component of the intermediate age
population of stars near the Galactic Center that does not exist in
Baade's Window. If we take
-6 as the upper limit to the AGB for the main
distribution of Galactic Center LPVs, then Fig. 19 of Vassiliadis
& Wood, with an assumed solar metal abundance, indicates a maximum
initial mass of
4
for this
group of stars. The empirical calibration of AGB tip luminosities from
LMC clusters (Frogel et al. 1990) suggests a similar or greater
mass.
As well as the main population of LPVs with
-6, there are several more luminous objects. The presence of stars
with an
-6 might be challenged by the
suggestion that these are in the foreground and much closer to us than
we have assumed. We noted in Sect. 1that at most a few of the
LWHN stars should be foreground objects. The following argument gives
strong evidence that at least one of these luminous stars is truely at
the centre or beyond. In trying to measure the angular size of the OH
masers of several of the LWHN stars,
van Langevelde et al. (1992) and
Frail et al. (1994) found that the size was unexpectedly large in
directions close to the Galactic Center. They concluded that
interstellar scattering in the region close the Galactic Center was
the most likely cause of the large observed size. One of the sources
observed to have a large size was LWHN 33 for which
Blommaert et al. (1997) derive
-6.2. In this situation, LWHN 33
must be behind the scattering region and close to the Center and it
cannot be a foreground object: its absolute bolometric magnitudes is
truely
-6.2. This result supports our conclusion
that other LWHN stars have a similar luminosity.
There are two LPVs of long period in our sample (LWHN 28, P
= 806 days; LWHN 116, P = 985) which have
-6.5, consistent with them being young AGB stars
with masses
(using the Vassiliadis & Wood
results). We also note that the two objects LWHN 70 and 90 are
exceptionally red and bright in L although they are relatively
faint in K so that periods could not be determined. However,
Van Langevelde et al. (1993) determined periods
1000 days for these two objects (the periods
belong to category 2, which are uncertain). Our photometry suggests
that LWHN 70 is not particularly luminous (
-3.1) but Blommaert et al. (1997) derive
= -6
suggesting this is an AGB star with
. LWHN 90 is
far too red for our bolometric correction to be applied. If we use our
colour and L magnitude with the
bolometric correction of Blommaert et al. (1997) for such a red star,
we derive
-7 as expected
for a massive AGB star with
. However,
Blommaert et al. (1997) derived a modest luminosity
=
-4.7. Two other stars with category 2 periods
1000 days are listed by Van Langevelde et al.
(1993). For one of these, LWHN 21, we derive an unambiguous period of
692 days (compared to 1733 derived by Van Langevelde et al.).
Blommaert et al. (1997) give a low bolometric luminosity of -3.7 for
the other one (LWHN 72). In summary, it appears that there are 2 to 5
massive (4-7
) AGB stars currently in the OH/IR
phase in the Galactic Center fields studied by LWHN. Since the
superwind or OH/IR phase of evolution in stars of
5
lasts about
105 years (Vassiliadis & Wood 1993;
Tanabé et al. 1997) while the corresponding main-sequence lifetime is
108 years, there should be about 1000
main-sequence stars of
5
near the Galactic Center for each OH/IR star with initial mass
5
. The existence of such
stars near the Galactic Center is consistent with a situation where
there is ongoing star formation occurring there.
A cautionary note regarding the quantitative estimates of AGB mass
is appropriate here. There are no theoretical computations of AGB tip
luminosities for stars with metallicity of 3.6
times solar, as estimated above for the Galactic Center LPVs. The
results of Vassiliadis & Wood (1993) suggest that AGB tip
luminosities for a given initial mass actually reduce with increased
metal abundance, which would suggest that the initial masses derived
above for AGB stars are underestimates. However, the theoretical
results are highly dependent on the mass loss scheme used and need to
be treated with caution. In fact, there are good reasons to believe
that the mass estimates of up to
7
given above are too high.
Firstly, in the LMC, the OH/IR stars with luminosity near the AGB
limit of
and probable masses of
5-7
attain periods of up to
1500 days (Wood et al. 1992) whereas the maximum
period attained near the Galactic Center is
1100
days (see the next section). Secondly, in the plane of our Galaxy, the
variable OH/IR stars attain pulsation periods of up to
2800 days (Herman & Habing 1985; Van
Langevelde 1990), this period presumably representing the most
advanced stage in the evolution of the most massive AGB stars
(
). Once again the maximum periods in the
Galactic Center region are less than in the Galactic plane. Although
the maximum mass for the Galactic Center AGB stars remains uncertain,
the high luminosities (see above) and long periods (see next section)
of these stars leave little doubt that the maximum mass is greater
than in the bulge.
© European Southern Observatory (ESO) 1998
Online publication: July 27, 1998
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