A&A 459, 849-857 (2006)
DOI: 10.1051/0004-6361:20065274
P. Koubský1 - P. Harmanec2,1 - S. Yang3 - M. Netolický4,1 - P. Skoda1 - M. Slechta1 - D. Korcáková1
1 - Astronomical Institute, Academy of Sciences,
251 65 Ondrejov, Czech Republic
2 -
Astronomical Institute of the Charles University,
V Holesovickách 2, 180 00 Praha 8, Czech Republic
3 -
Department of Physics and Astronomy, University of Victoria,
PO Box 3055 STN CSC, Victoria, B.C. V8W 3P6, Canada
4 -
Institue of Theretical Physics and Astrophysics, Faculty of Science, Masaryk University Brno, Kotlárská 2, 611 37 Brno, Czech Republic
Received 25 March 2006 / Accepted 28 April 2006
Abstract
Context. Binaries observed in the initial rapid phase of mass exchange between the components are very rare since the statistical probability of finding them is low. At the same time, thorough studies of them are extremely important for better understanding the process of large-scale mass exchange and possible mass loss from the system. One of these objects is probably Sgr.
Aims. By analyzing 35 new electronic spectra and numerous published spectral and photometric observations, we derived the new orbital elements, an upper limit to a secular period change, and also the peculiar RV curve of the blue-shifted H absorption. Possible models of the binary and its evolutionary stage are then discussed critically.
Methods. Reduction of new spectra was carried out with the IRAF and SPEFO programs. All orbital elements were derived with the FOTEL program and period searches were carried out using the phase-dispersion minimalization technique.
Results. The peculiar RV curve of the blue-shifted H absorption rules out the model of a coronal flow of matter from the brighter component. The presence of bipolar jets that are perpendicular to the orbital plane and similar to those found for
Lyr seems probable. An upper limit to the secular period decrease of 24 s per year was estimated and should be tested by future RV observations. The peculiar character of the line spectrum of the brighter component could also be understood as originating from a pseudo-photosphere of an optically thick disk rather than from a stellar spectrum.
Conclusions. Better understanding of the nature of Sgr will not be possible without interferometric resolution of the binary and especially without determining its orbital inclination.
Key words: stars: emission-line, Be - stars: binaries: close - stars: individual: Sagittarii
The emission-line binary Sgr (46 Sgr, HD 181615, HR 7342, HIP 95176, MWC 313;
V = 4
6 var.) is a very unusual object. Its radial-velocity
(RV hereafter) variations were discovered by Campbell (1899).
Wilson (1913) later demonstrated that
Sgr is a single-line spectroscopic
binary and derived its first orbital elements:
,
e=0.087,
,
and K=48.15 km s-1.
Duvignau et al. (1979) studied OAO3 and TD1 satellite UV spectra and concluded
that the far-UV spectra resemble a B0Ia object. From
the absence of measurable RV variations of high-excitation lines,
they deduced that the companion must be much more massive than
the object seen in the optical spectra. Parthasarathy et al. (1986) came to
similar conclusions on the basis of the IUE spectra.
The only real RV measurements of the fainter component by Dudley & Jeffery (1990)
are based on the cross-correlation of the far-UV spectra from
the short-wavelength camera on board the IUE satellite. These authors
also derived RVs of the component observed in the optical region, using all
high-dispersion SWP and LWP IUE spectra with a good S/N ratio. They
concluded that the eccentricity of the orbit is spurious and went on to derive
a circular orbit. Analyzing all the available RVs, they also
concluded that there were no detectable changes in the orbital period.
Perhaps the most important result of their study is the finding that
the mass ratio of the "invisible'' to visible binary component is 1.59,
i.e. that the invisible star is the more massive of the two. Dudley & Jeffery (1990)
also concluded that the luminosity ratio of the visible
to "invisible'' component must be about 100. However, in a later
paper Dudley & Jeffery (1993) show that such an extreme luminosity ratio is
inconsistent with the observed high flux shortwards of 1400
and that
the source of this excess is in the vicinity of the "invisible'' star.
Furthermore, the deconvolved spectrum of that component presented in
Dudley & Jeffery (1990) is characterized by a number of features of either
interstellar or circumstellar origin. The "invisible'' component might be
imbedded in an envelope, which also makes the RV determination and thus
the value of the system mass ratio uncertain. The detection of the secondary
RV curve therefore deserves an independent verification.
At the present stage of knowledge,
Sgr seems
to bear some similarity to
Lyr, because the component that
would normally be called primary
on the basis of its relative brightness could also be called
secondary on the basis of its mass.
To avoid confusion, we call the star dominating the optical
and UV line spectra star 1 in the rest of this paper and
its (probably more massive but less luminous) counterpart star 2.
A strong H emission in the spectrum of
Sgr was reported by
Campbell (1895) and by many other investigators. Maury (1925)
called attention to another peculiarity of the object: the optical
spectrum contains lines that resemble an F2 supergiant but also a B8 spectrum, and the relative strength of these two spectra varies in time,
the B8 spectrum being strongest at phases of a conjunction in which
star 2 is in front of star 1.
Such a variation in the relative strength of the hot and cool
spectra was not observed in later studies (based on better
spectra) by Greenstein & Adams (1947) and Greenstein (1950).
Even more puzzling is that the RVs of spectral lines corresponding to "F2'' and "B8'' spectra follow the same RV curve and are, therefore, associated with star 1
(Plaskett 1926). Maury (1925) also reports that another system
of absorption lines was sometimes observed on the Harvard spectrograms,
indicating RVs of about -280 km s-1. She pointed out that this system
of lines could hardly represent the orbital motion of star 2 since
no similarly high positive RVs had ever been observed. Bidelman (1949)
investigated these blue-shifted lines further, concluded that they
only occur around the phases of the conjunction with star 2 in front
of star 1 (but not every orbital cycle), and was probably the first
to suggest that they may indicate a gas stream from star 1 towards
star 2. Hack (1960) also analyzed high-dispersion H line
profiles and concluded that the blue-shifted absorption is indeed
strongest near the superior conjunction of star 1 but present with
progressively reduced strength throughout the whole orbital cycle.
Not considering her finding, Nariai (1967) collected and analyzed
available H
profiles, including new ones he obtained at Okayama,
and interpreted the blue-shifted lines as a supersonic flow from the stellar
corona of star 1 to a cone in the vicinity of star 2. Frame et al. (1995)
studied echelle spectra secured in the years 1988, 1991, and 1992.
They observed the blue-shifted H
absorption only on a single 1991
spectrum, although they observed the star near the superior conjunction
on seven other dates in 1988 and 1992. Instead, they detected strong
H
emission with a flat red wing, extending to some +500 km s-1.
Table 1:
Journal of RV data: the telescopes and spectrographs with which the RVs were obtained, denoted by the running numbers in column "Spg. No.'', are as follows:
1... James Lick 0.91-m refractor, Mills 3-prism spg, phot. plates;
2... Yerkes 1.02-m refractor, Bruce 1-prism spg. phot. plates;
3... DAO Plaskett 1.8-m reflector, Cass. IM or IS prism spgs., phot. plates;
4... Mt. Wilson Hooker 2.54-m reflector, coudé grating spg., 114-inch camera, phot. plates;
14... Mt Wilson Hooker 2.54-m reflector, coudé grating spg., 32-inch camera, phot. plates;
5... IUE satellite 0.45-m reflector, Cass. echelle spgs.,
SWP, LWP and LWR SEC Vidicon cameras;
6... CTIO 1-m reflector, Cass. spg., image tube & phot. plates;
7... MJUO McLellan 1-m reflector, Cass. fiber fed echelle spg., Reticon 1872;
8... SAAO 1.9-m reflector, Cass spg., image tube & intensified Reticon;
9... DAO McKellar 1.2-m reflector, coudé spg., SITe-4 4096
2048 CCD;
10... Ondrejov 2-m reflector, coudé spg., SITe-005 800
2000 CCD.
There is ample evidence that the continuum energy distribution
of Sgr is also composite and is a peculiar one from the IR to UV
(Jaschek et al. 1990; Burnashev 1981; Plavec 1986; Treffers et al. 1976; Hack 1980).
Gaposchkin (1945) obtained photographic photometry and concluded that
Sgr is an eclipsing binary with two low-amplitude minima, the eclipse
of star 2 being deeper than the opposite one. Eggen et al. (1950) published
the first photoelectric yellow and blue photometry from Lick. They also
obtained a double-wave light curve with 0
1 deep minima. However, they
pointed out that the minima occur some 12 days later than predicted
from the RV curve and noted that Gaposhkhin's minima also occur some
7 days after the spectroscopic prediction. Burnashev (1981), Malcolm & Bell (1986),
and Frame et al. (1995) also obtained photoelectric light curves and found
variations on a time scale of about 20 days but no clear evidence of
binary eclipses.
Hack & Pasinetti (1963) demonstrated that the chemical composition of star 1
is a peculiar one. They found a strong deficiency of hydrogen and
overabundancy of C, N, and some other ions. They interpreted it as
a consequence of a large-scale mass transfer between the binary components,
as did Plavec (1986). Dudley & Jeffery (1993) used published photometry,
spectrophotometry and line-blanketed hydrogen-deficient model
atmospheres to estimate the effective temperature of star 1 as
11 800
500 K. Leushin (2001) derived the effective temperature
of 13 500
150 K,
= 2
0.5, and the iron abundance for star 1 and
reviewed the previous studies of
Sgr done with his collaborators.
Schoenberner & Drilling (1983) have presented an evolutionary scenario for Sgr.
They suggest that star 1 is filling its Roche lobe and spilling its
now helium-rich envelope towards star 2 (i.e. that it undergoes the so called
case BB of mass exchange). Eggleton (2002) suggests that the binary had to
loose mass via envelope ejection, without removal of too much orbital momentum
so that a long orbital period and not too extreme mass ratio of 0.63 were
preserved.
Our observations consist of two series of red CCD spectra secured
at the Dominion Astrophysical Observatory (DAO hereafter) and in
Ondrejov. The initial reductions to 1D frames were carried
out in IRAF; by SY for the DAO and by MS for the
Ondrejov spectra. The final reduction, rectification, RV and
line-intensity measurements were carried out in SFEFO
(Horn et al. 1996; Skoda 1996) by PK and partly also by PH. The RV's were measured
interactively, comparing the direct and reverse images of the line profiles.
The zero point of the RV scale was corrected through the use of reliable
telluric lines. For the orbital RV
we finally used the mean of the strong, well-defined, and unblended
lines, which give consistent results: Ne I 6402.2455 and C II 6578.052
.
We also compiled all RVs available in the astronomical literature. In all cases when the heliocentric Julian dates (HJDs hereafter) were not tabulated in the original source, we used a computer program to derive HJDs for the data set in question. A few remarks on some of the data sources are in order.
File 5: IUE RVs Dudley & Jeffery (1990) tabulate Julian dates only as integers. We extracted HJDs for the mid-exposures for all their RVs and noted one misprint in their paper: JD of the spectrum LWR2276 is 2 443 757 (HJD 2 443 756.7729), not 2 443 557 as given in Table 1 of Dudley & Jeffery (1990). Note that these authors also tabulate 11 RVs of the secondary derived via a cross-correlation technique from some of the SWP spectra.
File 7: Mt. John echelle spectrograph Frame et al. (1995) tabulate
33 RVs in their Table 5, but a phase plot of these RVs gives a phase
diagram that differs slightly from the phase plot in their Fig. 6.
In particular, Fig. 6 does not show an RV value of -48 km s-1 observed
on JD 2 448 377.40 although the numbers of RVs listed in Table 5 and
plotted in Fig. 6 are the same. At our inquiry, Dr. Peter Cotrell kindly
informed us that the peculiar RV value, corresponding to an anomalous
H profile, was not used in their study. He also communicated another
observation missing from Table 5 but plotted in Fig. 6:
JD 2 447 464.50 with RV = +9 km s-1.
The journal of all RVs used in this study is in Table 1 and all individual RVs with HJDs and weights are tabulated in Table 2 (available only in electronic form).
All orbital solutions presented in this study were derived with the computer program FOTEL (Hadrava 2004). To derive new orbital elements, we proceed as follows.
Table 3:
Various orbital solutions for Sgr derived with the FOTEL program.
All epochs are in HJD-2 400 000 and the rms error is error per 1 observation
of unit weight. Systemic velocities of individual spectrographs are
identified by the running numbers corresponding to those
from Table 1. Rather large differences in the individual systemic
velocities probably do not reflect real changes but rather compensate
for the differences in line identifications and various subsets of lines
used by various authors to derive the published RVs.
We initially used only the RVs of star 1 and assumed a circular orbit.
First we derived a solution in which weights of all RVs were set equal to one.
This is solution 1 in Table 3.
Then, we made the weights of individual data sets inversely proportional
to the square of their rms errors per 1 observation,
which were derived in the preliminary solution with equal weights and
normalized in such a way that the mean rms per 1 observation from
this solution derived for all data was set equal to one.
All other exploratory solutions, as well as the final solution,
were derived with these weights (which are also given in Table 1).
Solution 2 is a circular-orbit solution for all RVs of star 1 with the
weights assigned to all individual data sets. This is the solution that
we adopt as the solution defining
the following linear ephemeris (1) to be used in this study:
In yet another solution 4, we fixed the value of the orbital period from ephemeris (1) and derived a circular solution for all star-1 RVs and for all RVs of star 2 obtained by Dudley & Jeffery (1990). Note that we tabulated the systemic velocity and the rms errors per 1 observation of unit weight separately for the primary and secondary RVs for this solution.
We also addressed the question of the secular constancy of the orbital period again. Using the RVs of star 1, we derived two circular-orbit solutions for older and more recent data. The main results are in Table 4.
This result is inconclusive. It indicates a 1- possibility
that the orbital period is secularly decreasing. We derived
a solution for all primary RVs allowing the period change as one
of the elements. We obtained a period change
dP/dt = (
d per day i.e. -24 s per year,
again a 1-
result. We only use this finding
as an estimate of an upper limit of the period change. Its verification with
new observations would be very important, however. If
Sgr is currently
in a phase of mass exchange, a decreasing orbital period would identify
it as an object in the rare initial rapid phase of mass exchange
before the role of components was interchanged
.
A more detailed inspection of Fig. 1 and Table 3 shows that the scatter of RVs around the mean orbital curve is higher than expected for easily measurable sharp lines of star 1. A part of the effect could be caused by the fact (mentioned earlier) that different lines seem to give slightly different RVs for the same spectrograms. To investigate the reality of effect, we therefore plot the new DAO and Ondrejov RVs vs. time in Fig. 2. One can see that there is little doubt about the reality of small cyclic variations on a time scale of some 20 days, i.e. similar to that of the observed light changes.
As already demonstrated by Bidelman (1949), Hack (1960),
Nariai (1967,1970), and Frame et al. (1995),
a blue-shifted H absorption line was observed on some orbital
cycles while missing in others. Published data on the presence and
absence of the blue-shifted absorption in H
were collected and schematically
represented by Nariai (1967,1970). Regrettably, there are
only a few published H
profiles from which one can measure central
intensity
of the H
absorption line. In particular,
we measured
in the H
profiles published by Hack (1960),
Nariai (1967), and Frame et al. (1995).
Table 4: Orbital solutions for two slightly overlapping subsets of older and more recent primary RVs to test the possibility of a secular change of the orbital period.
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Figure 1:
Upper panel: RV curves of both binary components
for solution 4. Bottom panel: the RV curve of the blue-shifted
H![]() ![]() |
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Figure 2: DAO RVs (filled circles) and Ondrejov RVs (empty circles) of star 1 plotted vs. time. Small systematic deviations from a smooth orbital RV curve are seen and data from two independent spectrographs prove their reality. |
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Figure 3:
Selected H![]() ![]() |
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All reliable observations suggest that in the cycles when the blue-shifted
absorption line is observed, it attains the maximum strength around
the conjuction with star 2 in front and virtually disappears at
the other conjunction. Note in particular that Frame et al. (1995) obtained
electronic H spectra of
Sgr covering parts of 12 consecutive orbital
cycles in 1988, 1991, and 1992. Absolutely no trace of the blue-shifted
H
absorption is seen on these high-S/N spectra (see their Fig. 8) with
the only exception of a single 1991 spectrogram
(phase 0.
345 according to ephemeris (1)),
which shows a strong H
absorption with an RV of -360 km s-1 and
a central intensity of 0.22. Instead, their H
profiles show
the H
emission with a steep blue edge and a very extended red
wing reaching to some +500 km s-1. The same type of H
emission
line with an extended red wing has already been observed by Greenstein (1943)
in July and September 1940.
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Figure 4:
Central intensities of the blue-shifted H![]() ![]() ![]() |
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In contrast, all H line profiles obtained at DAO and
in Ondrejov between August and October 2005 do show the
blue-shifted H
absorption of variable intensity; see
a representative selection of the H
profiles in Fig. 3.
The radial velocity of blue-shifted H
absorption was also derived
by comparing direct and reverse images of the line profile.
Our measurements refer to the line core. The values of RVs
from the phases when the satellite profiles are very faint are of course
less reliable. On the other hand, the RV curve we obtained is smooth,
changing continuously from phases where the satellite component dominates
to phases where it gets faint. This could hardly be so
if blending with some photospheric line takes place.
In the interval of phases in question (around 0.8), the photospheric lines
move in antiphase to the RV curve of the blue-shifted
H
absorption line.
The RV curve of the blue-shifted H
absorption is shown in the bottom
panel of Fig. 1. From it we conclude that the RV behaviour
of the H
absorption does not support the Nariai (1967) interpretation in
terms of a supersonic flow from the corona of star 1 extending to the vicinity
of star 2. Very regrettably, our observations do not cover
phase 0
25 of the superior conjunction of star 1 where one would expect
the largest negative RVs for his model. However, if the absorption
were to originate from a stream flowing between the two stars, there would be
no absorption seen in the elongation with star 1 approaching. Even if there
were some self-absorption in the flow, its RV at these phases would be
close to the systemic velocity, which is certainly not the case.
We do confirm, however, the Nariai (1967) finding that the line attains
its maximum strength at the superior conjunction of star 1
and minimum strength at the opposite one, when star 1 is
in front - see Fig. 4. Note also a remarkable phase
coherence in the line-strength variations for the earlier observed and new
data. In other words, if the blue-shifted absorption is present in
the spectra, its strength at any given orbital phase remains the same over
many orbital cycles. This seems to indicate that the alternation between
the presence and absence of this feature is probably a geometric, not
physical effect.
Several authors have claimed that the blue-shifted absorption lines are
also present in other Balmer lines and in some helium lines.
In Fig. 5 we show a montage of
the Ondrejov profiles of H,
H
,
H
,
and
He I 6678. It is seen that, in the phases when the blue-shifted H
absorption is strong, it is also observable in H
and H
.
There is no evidence, however, for a blue-shifted satellite line in He I 6678.
Plaskett (1926) published RVs of several putative blue-shifted
absorption components of H,
H
,
H
,
and even H
measured throughout the whole orbital cycle.
In Fig. 6 we show the RV of one H
satellite line transformed
to the RV scale of Fig. 1. The resulting RV curve is reminiscent
of the RV curve of star 1. This corroborates the conclusion by
Greenstein (1950) that Plaskett (1926) probably misidentified
some fainter lines of star 1 for the blue-shifted components of
the Balmer lines.
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Figure 5:
Blue-shifted absorption lines of several Balmer lines from
the Ondrejov CCD spectrograms. The same blue shift is observed
for all three Balmer lines from similar orbital phases.
Note also a good coherence of the phase
variations of H![]() ![]() |
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Figure 6:
The RV curve of a blue-shifted H![]() |
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Figure 7:
Orbital RV curve of the H![]() |
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Figure 8:
A comparison of two Ondrejov H![]() |
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To have some idea about the location of the H emission within the binary
system, we measured the RV of the H
emission peak. Its variations with
orbital phase are shown in Fig. 7. There is some disturbance
near the conjuction with star 1 in front of star 2, leading to negative values
of RV. Figure 8 compares two Ondrejov spectrograms
obtained in orbital phases 0
127 and 0
753. One can see that near
the conjunction the emission line gets much wider, and there is probably
some additional source of emission. Outside the conjunction phase, however,
the emission RV seems to follow a sinusoidal variation almost in phase with
star 1 but with a reduced amplitude. A formal orbital solution with the
data near phase 0
75 omitted leads to
K = 10.1
5.2 km s-1,
= 3.4
0.8 km s-1,
= HJD 2 433 009
2,
the rms per 1 observation being only 2.9 km s-1.
A similar result was also found for the Fe II 6432
emission line,
except that its RV curve is not disturbed near phase 0
75.
This seems to indicate that the bulk of the emission originates in some
area located on the same side from the binary centre of gravity as star 1.
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Figure 9: Several orbital light curves for the V-magnitude observations found in the astronomical literature. |
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With the new accurate ephemeris, we compiled several representative
V magnitude photometric data sets from the astronomical literature
and plotted them vs. phase in Fig. 9. In particular we used
the photometry of Eggen et al. (1950), Burnashev (1981), Manfroid et al. (1991),
Sterken et al. (1993), Malcolm & Bell (1986), Frame et al. (1995), and Hipparcos
photometry (Perryman & ESA 1997) converted to the Johnson V magnitude
following Harmanec (1998). One immediately sees that there is no evidence of
photometric eclipses or ellipsoidal variations, although the observations
of the lowest brightness of the object seem to cluster near phases 0
2-0
3
and 0
7-0
8, which may explain the early reports of eclipses.
Figure 9 also illustrates the conclusion of Frame et al. (1995) that
the amplitude of variations varies secularly with time. This implies
that one remains empty-handed as far as the knowledge of the orbital
inclination, and therefore binary masses, is concerned.
Analyzing the individual data sets for periodicity, we found that
a time scale of 35-40 days is usually detected, leading to a double-wave
light curve with two minima of different widths. Note that
Malcolm & Bell (1986) found a twice shorter periodicity for a single-wave
curve so that their findings probably refer to the same time
scale of variations as ours. For illustration, we show
a phase plot of their V photometry for the best period of
in Fig. 10.
We intend to re-observe all comparison stars used by different observers to be able to homogenize the existing data sets and combine them for a future period analysis.
After a critical re-axamination of the observational evidence, we
conclude that one should remain open-minded when interpreting
the geometry and evolutionary stage of Sgr. We also suggest
a few hints.
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Figure 10:
A phase plot of Malcolm & Bell (1986) V photometry
for the best "period'' of
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Table 5:
Probable basic physical properties of Sgr based on
orbital solution 4 of Table 3 derived for three possible values of
the orbital inclination. Tabulated are masses M1 and M2, the
semi-major axis A, Roche-lobe radii
and
,
and the minimum and maximum angular dimensions of the semi-major axis
and
derived for the range given by the Hipparcos parallax and its error
.
Nevertheless, it seems that the most powerful of the already operating optical
interferometers might be potentially able to resolve this binary,
an attempt that is worth the effort. At the same time it is clear that
the interferometry of Sgr could substantially help to restrict the possible
classes of models for this peculiar binary.
The most important would be an inteferometric determination of
the orbital inclination of
Sgr which - together with an independent
confirmation of the RV curve of star 2 - would finally settle
the question of the true masses of both bodies.
Acknowledgements
We acknowledge the use of the programs SPEFO and FOTEL, made available by their authors Drs. Jirí Horn and Petr Hadrava. Our thanks are also due to Dr. Peter Cottrell, who kindly clarified the problem with Mt. John's RVs for us. Very relevant comments, suggestions, and a constructive criticism by the anonymous referee helped to improve the paper and are gratefully acknowledged. We profitted from the use of the electronic bibliography maintained by the NASA/ADS system. Czech authors were supported by the research plan AV 0Z1 003909, from project K2043105 of the Academy of Sciences of the Czech Republic, and from the grants GA CR 205/02/0788 and 205/06/0584. The research of PH was also supported from the grant GA CR 205/06/0304 of the Czech Science Foundation.