A&A 396, 589-598 (2002)
DOI: 10.1051/0004-6361:20021416
P. Gil-Pons1 - E. García-Berro1,2
1 - Departament de Física Aplicada, Universitat
Politècnica de Catalunya, c/Jordi Girona s/n, Módul
B-4, Campus Nord, 08034 Barcelona, Spain
2 -
Institute for Space Studies of Catalonia, c/Gran
Capitá 2-4, Edif. Nexus 104, 08034 Barcelona, Spain
Received 25 March 2002 / Accepted 10 September 2002
Abstract
The evolution of a star of initial mass 9 ,
and metallicity
Z = 0.02 in a Close Binary System (CBS) is followed in the presence of
different mass companions in order to study their influence on the final
evolutionary stages and, in particular, on the structure and composition
of the remnant components. In order to do that, we study two extreme
cases. In the first one the mass of the secondary is 8
,
whereas in the second one the mass was assumed to be 1
.
For
the first of those cases we have also explored the possible outcomes of
both conservative and non-conservative mass-loss episodes. During the
first mass transfer episode, several differences arise between the
models. The system with the more extreme mass ratio (q=0.1) is not
able to survive the first Roche lobe overflow (RLOF) as a binary, but
instead, spiral-in of the secondary onto the envelope of the primary
star is most likely. The system formed by two stars of comparable mass
undergoes two mass transfer episodes in which the primary is the donor.
We have performed two sets of calculations corresponding to this case in
order to account for conservative and non-conservative mass transfer
during the first mass loss episode. One of our main results is that for
the non-conservative case the secondary becomes a Super-AGB star in a
binary system. Such a star undergoes a final dredge-up episode, similar
to that of a single star of comparable mass. The primary components do
not undergo a Super-AGB phase, but instead a carbon-oxygen white dwarf
is formed in both cases (conservative and non-conservative), before
reversal mass transfer occurs. However, given the extreme mass ratios
at this stage between the components of the binary system, especially
for the conservative case, the possibility of merger episodes remains
likely. We also discuss the presumable final outcomes of the system and
possible observational counterparts.
Key words: stars: evolution - stars: binaries: general - stars: AGB and post-AGB - stars: white dwarfs
This is the second of a series of papers in which we intend to explore
extensively the evolution of intermediate mass close binary systems
(IMCBS). IMCBS are defined as those systems in which the primary
component develops a partially degenerate carbon-oxygen core, after
burning central helium in non-degenerate conditions. In particular, we
focus on the evolution of heavy-weight intermediate mass stars, which
have primary masses between
and
.
These
stars are thought to ultimately develop ONe cores (Ritossa et al. 1996), and populate the brightest portion of the
Asymptotic Giant Branch. In this paper we follow the evolution of both
components of the binary system from their main sequence phase until
late evolutionary stages, paying special attention to the carbon burning
phase, and to the final objects we can encounter afterwards.
Even though the evolution of isolated intermediate mass stars has
recently been analyzed - see, for instance, Ritossa et al. (1999), and references therein - the evolution of this kind of
star in binary systems has been very little studied. Probably one of
the reasons for the absence of this kind of model in the literature is
that the conversion of a CO core into an ONe core involves following the
evolution of an unstable nuclear burning flame as it propagates first
towards the center of the star and then outwards. The calculation is a
very delicate and time-consuming task since the propagation of the flame
front is not steady, but is interrupted by a series of rather violent
shell flashes (García-Berro et al. 1997). The flashes
are violent because the nuclear reaction rates are very temperature
sensitive and the material is partially degenerate. Moreover, because
of the near discontinuity in the physical variables at the flame front,
following the inward motion of the flame requires very good spatial
resolution and short time-steps. Furthermore, the evolution of stars
within this range of masses in binary systems is even more complicated,
since the presence of a close companion can dramatically alter the
evolution of the primary star and its final outcome. As a result the
only recent calculation in which a heavy-weight intermediate mass star
is followed from its main sequence phase until carbon is exhausted in
the core is that of Gil-Pons & García-Berro (2001), who followed
the evolution of a
star with solar metallicity in a binary
system. Moreover, most of the developments in the field of binaries
frequently disregarded the study of AGB stars, arguing that this phase
is prematurely quenched in binary systems, due to significant mass loss
in the evolution previous to the AGB phase. Important exceptions are
the works of Jorissen (1999) and Smith et al. (1996) who proposed the
existence of AGB stars in binary systems and, more recently, Van Eck et al. (2001), who found observational evidence of three AGB stars in
binary systems.
In our previous paper (Gil-Pons & García-Berro 2001), we
proposed a scenario in which a
star could evolve to
become a Super-AGB star and, ultimately, an ONe white dwarf, in spite of
the fact that all the hydrogen-rich envelope and most of the helium
layer had been lost in previous phases of the evolution. In this paper
we consider again such an scenario but we relax some of the simplifying
assumptions that were made there. In particular, we compute the
evolution of a
primary (the initially most massive star)
of the binary system, and we explore the effect of the mass of its
companion. In doing so, we study two extreme cases; in the first one
the initial mass ratio (q0) is close to one (scenario 1), whereas in
our second calculation this ratio is less than 0.2 (scenario 2). In
fact, q0=0.2 is a theoretical limit under which spiral-in of the less
massive star onto the most massive and finally, the merger of
components, cannot be avoided - see, e.g., Vanbeveren et al. (1998a).
Han et al. (2000) did a comprehensive study of the evolution of IMCBS
in order to determine the influence of the initial mass ratio and
orbital period on the final parameters of close binary systems whose
primary initial mass is in the range
to
.
Although they only consider the conservative case, a significant
variation in the primary remnant masses is observed. In this paper we
focus on more massive primaries - which have been little studied -
and we extend our previous calculations in order to consider the
evolution of the secondary. Furthermore, since there are still many
uncertainties in relation to whether mass transfer is conservative or
not, we have computed two evolutionary sequences for the first of our
scenarios, and we have considered the two extreme cases that may occur
during the first mass loss episode:
The paper is organized as follows. In Sect. 2, we present in detail our evolutionary scenarios. Section 3 is devoted to the study of the evolution of the first of the above mentioned scenarios in which a massive companion is assumed, whereas in Sect. 4 we explain in detail the evolution in the case of a low-mass secondary. The phase of reversal mass transfer in the first of our scenarios is fully explained in Sect. 5. Finally in Sect. 6 we discuss and summarize our major findings.
A general overview of the evolutionary scenarios mentioned in Sect. 1 can be found in Fig. 1, for the cases 1.a and 1.b, and in Fig. 2 for case 2.
The cases 1.a and 1.b (see Fig. 1) undergo two mass loss episodes. In
both cases, the primary remnant has lost most of its hydrogen-rich
envelope after the first one, but for the case 1.a, this mass is
,
slightly higher than the mass of the case 1.b remnant
(
). Depending on whether the mass transfer
process is conservative or not the final mass of the secondary changes
accordingly, with the corresponding values 8.7 and
for the
and
cases, respectively. The final
orbital periods also change after this phase, to 200 and 120 days
for the conservative and non-conservative case respectively.
Furthermore, using the prescription of Iben & Tutukov (1985) for
non-conservative mass loss, we assume that that the angular momentum
losses during the common envelope phase are as high as 70-80%.
Orbital shrinkage is a direct consequence of this kind of mass loss, and
therefore one can expect that the orbit of the binary system will be
wider for the case of conservative mass transfer. We will come back to
this point in Sect. 3.2. Despite these differences, the next mass transfer
episode for both cases takes place during the helium-shell burning
phase. At the end of the second mass transfer episode, the remnant of
the primary consists of a carbon-oxygen core of about
surrounded by a helium envelope of about
.
We would like to emphasize that considering convective
overshooting might have led to somewhat more massive remnants for the
primary. However, as we intend to keep consistent with our previous
paper, and because no significant differences are expected, we have
adopted the standard Schwarzschild criterion with no overshoot.
![]() |
Figure 1: General overview of the evolution of the binary system and outline of the possible final outcomes, for the cases 1.a and 1.b, see text for details. CHeB stands for central He-burning, SHeB for shell He-burning, SAGB for Super-AGB and the rest of the acronyms have their usual meaning. |
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At the end of the second mass transfer processes, the masses of the
companions are 9.3 and
and the orbital periods are
1000 and 700 days for cases 1.a and 1.b, respectively. After a very
short phase in which carbon is partially burnt off-center in the core of
the primary, the bulk of the carbon burning phase is ultimately avoided,
and the remnant of the primary star is a CO white dwarf. At some point
during the white dwarf cooling phase, the secondary becomes a red giant
and reversal mass transfer ensues. The final outcome depends on the
mass ratio of the binary system. Given that at this point the mass
ratio is
,
spiral-in of the white dwarf onto the secondary
and, consequently, a merger of the two stars cannot be discarded. If,
otherwise, the system is able to survive in spite of the extreme mass
ratio, one can expect several different outcomes, which will be
discussed in detail in Sect. 5.
Figure 2 shows schematically the main evolutionary stages of our third
evolutionary scenario. For this case there is only one episode of mass
transfer, occuring when the primary reaches red giant dimensions. The
most likely outcome is spiral-in of the secondary onto the primary and,
ultimately, the merging of both components. This happens as a
consequence of their initial mass ratio (
). We simulate
the process as a thermohaline mixing - see, for instance, Vanbeveren
et al. (1998b). That is, we assume that all the material belonging to
the secondary mixes instantaneously and homogeneously with the
hydrogen-rich envelope of the primary. The final result of this process
is a
star. However, this star will have the core of a
star. A consequence of the merger is that the
composition of the envelope is slightly changed, showing a slight
overabuncances of C and O and an underabundance of N. After the merger,
the product continues its evolution as a single star, very much in the
way described by García-Berro & Iben (1994).
![]() |
Figure 2: General overview of the evolution of the binary system, for the case 2 - see text for details. |
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In Fig. 3 the evolutionary tracks in the Hertzsprung-Russell diagram of the primary components of the CBS for both the case 1.a - left panel - and for the case 1.b - right panel - are shown. The mass loss episodes are represented as dotted lines. The primary starts losing mass just after the hydrogen core burning phase for case 1.a, whereas for the case 1.b the primary begins its RLOF when already climbing the red giant branch (RGB). Once this process is completed for both stars, they make their way to the left of the Hertzsprung-Russell diagram and become blue stragglers, since they are almost completely deprived of their hydrogen-rich envelopes. During this phase the bulk of core He-burning occurs, and once helium is depleted in the inner core, the primaries evolve towards decreasing effective temperatures.
![]() |
Figure 3: Evolutionary tracks of the primary components of cases 1.a and 1.b in the Hertzsprung-Russell diagram. Mass transfer episodes are shown as dotted lines. |
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The second mass loss episode for both cases starts when helium burns in
a shell. At this point the luminosity increases at almost constant
effective temperature for both cases. The mass loss processes in both
cases continue while the primaries climb the giant branch and finish
when the evolutionary paths of these stars are well advanced on the
horizontal track for case 1.a, and just at the beginning of the
horizontal track for case 1.b. Finally, once the primary components of
the CBS are deprived of most of their envelope they become white dwarfs
and follow the corresponding cooling tracks. On the contrary, the two
secondary components (Fig. 4) look very different from each other. In
particular, the narrow loop feature that appears in both cases before
the top of the main sequence is reached are a consequence of the very
high accretion rates at which mass transfer proceeds during the first
RLOF. The second phase of accretion is less apparent, as it takes place
at much smaller accretion rates. Once both episodes of mass transfer
are over, the secondary for case 1.b evolves rather regularly, following
the path on the Hertzsprung-Russell diagram of an isolated
,
which is very similar to that of a
star studied
by García-Berro et al. (1997).
The evolution of the secondary for the case 1.a requires a more careful
explanation. It starts its evolution as a normal
single
star. However, RLOF from its companion drives a mass transfer process
during its early core hydrogen burning phase that allows it to almost
double its mass (up to
)
in a relatively brief
interval of time. This translates into an important increase in the
luminosity (see the left panel of Fig. 4) and a slight increase in the
effective temperature. Accordingly, the surface radius of the secondary
increases considerably (from
to about
)
but not enough to overflow its own Roche lobe. This effect
has been already noted by Vanbeveren et al. (1998a) as being a direct
consequence of the fact that the thermodynamic structure of the
secondary is not able to adapt to the rapidly accreted matter.
The second mass transfer episode also happens during the core hydrogen
burning phase of the secondary. During this episode the amount of
accreted mass is relatively modest (
), and the
accretion rates are much lower. The phase of core hydrogen burning for
the secondary finishes shortly after this second accretion process. The
subsequent evolution in the Hertzsprung-Russell diagram is similar to
that of a high mass single star.
![]() |
Figure 4: Evolutionary tracks of the secondary components of cases 1.a and 1.b in the Hertzsprung-Russell. Mass transfer episodes are shown as dotted lines. |
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As already mentioned, mass transfer in case 1.a begins shortly after the
primary burns hydrogen in a shell. The presence of a radiative envelope
surrounding the mass donor allows for moderate mass transfer rates. In fact, one can easily distinguish two phases in the mass loss episode -
see the bottom panel of Fig. 5. The first phase corresponds to the
beginning of mass transfer, proceeds at thermal time scales and the
typical mass loss rates are
.
This phase lasts for about 104 yr, and the bulk of the
hydrogen-rich envelope is lost during it. The second phase takes place
at much lower mass loss rates (
-
), but it is considerably longer, lasting for,
approximately,
yr. The upper panel of Fig. 5 shows
how the surface abundances change as the mass transfer proceeds. The
begining and the end of the mass transfer episode are clearly marked by
the thin vertical lines. As it can be seen in this panel, at the end of
the first mass loss episode the only CNO ash in the remaining envelope
is 14N.
For case 1.b the primary starts losing mass while it is climbing the red
giant branch. Consequently, the presence of a deep convective envelope
during this evolutionary stage induces mass loss to proceed on fast time
scales, during a common envelope phase, as was the case studied in
Gil-Pons & García-Berro (2001). Since the first RLOF takes place
on very fast time scales, and since the surface radius of the primary is
well over the value of its Roche lobe radius during most of the time, we
let the mass transfer proceed in such a way that the Roche lobe radius
(Eggleton 1983) is proportional to the surface radius. As we intend to
study very non-conservative mass transfer episodes, we assume that the
Eddington accretion limit is valid, and we do not allow for
super-Eddington accretion rates, as was done in the calculations of
Podsialdowski et al. (2001) and Portegies-Zwart et al. (2000). That
is, we consider that all the matter lost by the primary at a rate higher
than the Eddington limit is lost by the system. During most of the
first RLOF episode, mass loss proceeds at rates larger than
and we obtain that only about
of the
mass lost by the primary is actually accreted by the secondary.
![]() |
Figure 5:
Upper panel: evolution of the surface abundances of the
primary for the case of a ![]() |
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As for the conservative case, for this one we can also distinguish a
fast and a slow phase in the mass transfer episode - see the bottom
panel of Fig. 6. However, in this case the fast phase yields much
larger mass loss rates (
). The whole RLOF episode lasts for about
yr.
Again, in the top panel of Fig. 6 the evolution of the surface
abundances as the mass transfer episode proceeds is shown. It is easy
to see that the final composition at the surface is very similar in both
case 1.a and in case 1.b. Our final hydrogen abundance at the end of
the first mass loss episode turns out to be
0.3 in both cases,
in very good agreement with the results of the calculations of de Loore
& Vanbeveren et al. (1995). The second mass loss episode proceeds at
much lower transfer rates, and thus the secondary is able to accrete all
the matter that is lost by its companion, with little modification of
its thermodynamical magnitudes (see the right panels of Figs. 3 and 4). This process starts after the development of a convective envelope
surrounding the primary as a consequence of helium shell burning and it
lasts until carbon is partially burnt in the helium-exhausted region.
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Figure 6:
Upper panel: evolution of the surface abundances of the
primary for the case of a ![]() |
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Figure 7 shows the abundance profile of the primary for our case 1.a, and
in Fig. 8 we show those abundance profiles for the case 1.b. The
differences of composition for the central regions are negligible, the
most abundant nuclei being oxygen ( = 0.62), and carbon
(
= 0.37) in both cases. Minor differences in the shapes of
the chemical profiles of the two cores appear as a consequence of the
peculiarities of their evolution, specially in those phases in which
mass loss or convection play an important role. The main difference is
the size of the core, that happens to be smaller for the case 1.b
(
), when compared to that of the case 1.a
(
). This is a direct consequence of the
fact that, because the orbit is wider for case 1.a, the corresponding CO
core has more time to grow before a new RLOF occurs. The similarity
between the primary cores before reversal mass transfer for the two
cases supports one of the main hypotheses in Gil-Pons &
García-Berro (2001) at this stage of the evolution, namely that
the core of the primary is only slightly affected by the initial mass of
the secondary.
The secondary for case 1.a undergoes two accretion episodes while still burning hydrogen in its core. Computationally, we have treated accretion in the usual way, the so-called standard accretion model. That is, we have assumed that the infalling matter has the same specific entropy as the surface layers of the mass-gaining component - see, for instance, Pols (1994). As noted by Vanbeveren et al. (1998a), using this approximation has two important consequences. The first one is the sudden increase in the luminosity of the mass-gaining component which has been already commented on Sect. 3.1. The second one is the increase of the convective core as accretion occurs. This effect is clearly illustrated in the top panel of Fig. 9. The times at which the two accretion episodes begin are clearly marked by the thin vertical lines. The total mass of the star is also shown as a thick solid line. Since the second episode occurs at a much more modest pace, this effect is not as evident as it was for the first one. On the other hand, the change in the surface composition of the secondary is more evident in this case, as the accreted material onto its hydrogen-rich envelope is almost pure helium. The change in the composition has the effect of driving the evolution in the Hertzsprung-Russell diagram to higher luminosities than those corresponding to a star of the same mass and a non-polluted envelope.
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Figure 7: Abundance profiles of the primary at the end of He burning for case 1.a. |
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After these two accretion episodes, the secondary evolves as a single
star. When the secondary reaches red giant dimensions,
superwinds start playing an important role. We have used the
prescription of Jura (1987) in order to take into account the mass-loss
rates due to stellar winds (see also Vanbeveren et al. 1998b). For our
case the mass-loss rate turns out to be
.
The effect of mass loss from the secondary is
clearly seen for long enough times in the top panel of Fig. 9.
However, it should be taken into account that several uncertainties
surround this evolutionary phase, which can significantly affect the
evolution of the binary system beyond this point. The main source of
these uncertainties is the fact that the radius of the secondary gets
very close to that of its Roche lobe. Therefore, it is difficult to
ascertain whether or not reversal mass transfer occurs. Furthermore,
there is the possibility of enhancement of the wind mass-loss rates due
to the presence of a close companion or, even, partial capture of the
wind by the primary remnant, as happens in symbiotic systems (Iben &
Tutukov 1996). Moreover, when carbon is ignited at the center of the
secondary, this star has not yet abandoned the red giant branch, and
both the helium and the hydrogen burning shells remain active. We have
not followed in detail the carbon burning phase because of these
uncertainties. However, at this point of the evolution the resultant
core is made of CO, with an oxygen abundance of
and a
carbon abundance of
.
The size of this
helium-exhausted core is
and it is surrounded by a
helium envelope. However, since the uncertainties in the wind mass-loss
rates are significant, we intentionally refrain from giving a precise
value of its mass.
![]() |
Figure 8: Abundance profiles of the primary at the end of He burning for case 1.b. |
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Figure 9: Convective regions for the evolution previous to carbon burning, for the secondaries of both our case 1.a (top panel) and 1.b (bottom panel). |
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Figure 10: Evolution of the main physical parameters of the secondary during carbon burning for case 1.b. |
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The evolution of the secondary component for case 1.b (non-conservative)
mass transfer presents some differences with respect to that of the case 1.a. These differences are not important, as far as the two accretion
processes are concerned: in fact, both of them occur during the core
hydrogen-burning phase of the mass gainer, and are accompanied by the
same effects of overluminosity and extra growth of the helium core (see
the top panel of Fig. 9). Note however that the first mass transfer
episode starts earlier for case 1.a than for the case 1.b, as can be
seen in Fig. 9. The differences become more remarkable once the
accretion process is finished. The main reason for these differences
is, of course, that the final masses of the secondaries are very unlike:
for the case 1.a, and
for the case 1.b, after the two accretion episodes. The new
component evolves as an intermediate mass star, very much in the same
way as the isolated
star studied by García-Berro
et al. (1997). In spite of the important angular momentum loss that
the system undergoes during the first mass transfer process, the final
orbital separation after the second episode is such that the secondary
component does not fill its Roche lobe until it reaches the Super-AGB
phase. In fact, this star is able to undergo a dredge-up episode and a
series of thermonuclear flashes caused by carbon burning in a degenerate
core (see the top and bottom panels of Fig. 10).
Opposite to what happens to the
model studied by
García-Berro et al. (1997), this dredge-up episode is not a
consequence of energy release during the first carbon burning flash.
Instead, it is the gravothermal energy release during the core
contraction phase which follows the exhaustion of central helium that is
responsible for the expansion and cooling of the outer envelope.
Figures 10 and 11 help to explain this phenomenon by showing the
evolution of the most important model parameters during this dredge-up
episode. In particular the upper panel of Fig. 10 shows the evolution
of carbon and helium luminosity,
and
,
respectively. In this panel the flashes due to carbon ignition in a
degenerate medium are apparent for times longer than
.
As usual, these thermonuclear flashes are
accompanied by the formation of associated convective regions in the
inner core (see the bottom panel of Fig. 10), and by expansion and
cooling of the upper layers, leading to the temporary extinction of the
helium burning shell. Note, however, that the dredge-up episode is not
a consequence of the carbon flashes, since these flashes occur well
after the fast penetration of the convective envelope, which occurs at
.
Instead, in the lower panel
of Fig. 10 it can be seen that the time at which the fast inner
penetration of the base of the convective envelope is taking place is
practically coincident with the time at which the surface radius
increases at an almost constant surface luminosity, owing to an increase
of the core density. In order to reinforce this argument, in Fig. 11
we show several relevant physical quantities for three selected models
at three different times: before the fast advance inwards of the
convective envelope (model a), during the advance (model b),
and after it (model c). The location of these models is clearly
shown in Fig. 10 and their corresponding times are labeled in Fig. 11.
The upper panel of Fig. 11 shows the terms involved in the equation of
energy conservation for model a. The minimum in the solid line
corresponds to the gravitational energy release that, as can be seen, is
transformed into work of expansion (dashed line), rather than into
internal energy (dotted line). Note as well that this is done very
close to the inner edge of the contracting envelope. The work of
expansion induces a considerable increase in the surface radius of the
star (central panel of Fig. 10). Ultimately, the cooling of the
external layers has, as a main consequence, the increase in the envelope
opacity - see the middle panel of Fig. 11. Consequently the
radiative gradient increases and, at a certain point, allows the
convective envelope to move inwards.
![]() |
Figure 11: Energy variation rates for model a - upper panel - and opacity and temperature gradient profiles for models a, b and c - middle and bottom panels, respectively. |
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Figure 12: Abundance profiles of the secondary component at the end of the calculation for case 1.b. |
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Finally, Fig. 12 shows the ONe core of the secondary component for the
case 1.b. The size of this core is of about
,
and the
main isotopes are oxygen (
), and neon (
). As can be seen in this figure, the neon-rich core is
surrounded by a CO rich layer of irregular profile, as well as by a
helium rich envelope, in agreement with the results of Gil-Pons &
García-Berro (2001).
At the end of the carbon burning phase the surface radius has increased
considerably, reaching values close to the Roche lobe radius (
). At this evolutionary stage the situation is similar
to that found in Gil-Pons & García-Berro (2001). Thus, it is
difficult to determine the mass-loss rates, as there is the combined
effect of RLOF and the typical superwinds of the AGB phase.
Perhaps the main uncertainty in this phase of the evolution is the fact
that stellar winds from such a massive secondary might be enhanced -
see, for instance, Iben & Tutukov (1996) - by the presence of a
companion. Another uncertainty has to do with the efficiency of the
secondary component in capturing the stellar wind. In order to provide
an estimate of this efficiency, we proceed as in Iben & Tutukov (1996)
and we assume that the remnant of the primary is able to accrete all the
wind flowing into a disk that has the Bondi-Hoyle radius
![]() |
(1) |
where
is the mass of the remnant of the primary,
and
is the velocity of the wind flowing from the secondary,
that we have taken to be about 25 km s-1, which is a typical value for red
giants (Iben & Tutukov 1996). Note, however, that this value can vary
significantly for different red giant stars. Under this assumption the
accretion rate is:
![]() |
(2) |
where
is given by the expression of Jura
(1987), and A is the orbital separation. We get that the accretion
rates are between
and
of the wind mass loss rates.
Therefore,
-
.
Analyzing this mass transfer process is certainly a very
delicate task, as the extreme initial mass ratio between components is
such that the possibility that the system would undergo a merger episode
is non negligible. Hence, hydrodynamical simulations should be carried
out in order to account for it in a realistic way - see, for instance,
Hjellming & Taam (1991), Iben & Livio (1993), Lai et al. (1993) and,
more recently, Taam & Sandquist (2000), and Ivanova & Podsiadlowski
(2001). These calculations are beyond the scope of this paper, but,
still, we will make some reasonable estimates of the possible final
outcomes in Sect. 5.
The system composed of a
and a
star
evolves in the following way: because the more massive star evolves
much faster than its companion, it undergoes RLOF while the secondary
component is still burning hydrogen in its core. The main
characteristic that determines the evolution of the system is, in this
case, the initial mass ratio between components, q = 0.1. This value
is smaller than the critical mass ratio and, thus, no matter whether or
not the initial period is such that the primary component is able to
develop a deep convective envelope, the system is forced to undergo a
merger episode (Vanbeveren et al. 1998a). This merger occurs at time
,
and in order to follow it we use
the thermohaline mixing approximation. Once the merger has ended, the
resulting single
star behaves very much like the
star studied by García-Berro & Iben (1994).
We have followed the evolution of a 9
star with solar
metallicity in close binary systems with different initial orbital
parameters. Our main goal has been to analyze the influence of the mass
of the secondary component on the final possible outcomes of the close
binary systems. In particular, we have computed a binary evolutionary
scenario that allows the formation of Super-AGB stars in close binary
systems. Our analysis encompassed all the relevant evolutionary phases,
starting from the main sequence of both components, following the helium
burning phase, and, if necessary the carbon burning phase. In summary,
we have studied the following three cases:
For the last of these three cases the first mass transfer episode leads
to the merger of components, due to the initial mass ratio between
components: .
The resulting
single
star behaves very much like the normal 10
star studied by
García-Berro and Iben (1994), once the merger episode is over.
Cases 1.a and 1.b are more complicated. These two binary systems
undergo two mass transfer episodes. This kind of evolution leads to
very similar primary remnants: both of them are massive CO white
dwarfs,
.
However, the two systems
are very different in orbital period and masses of the secondaries. For
case 1.a the system is composed of a primary remnant and a
main sequence star, with a hydrogen-rich envelope highly
polluted by the CNO products from the primary remnant, dredged-up by
convection and expelled in its second RLOF, and also by the products of
its own helium burning shell. The secondary for case 1.b is a
star. In many aspects the evolution of this star resembles
very much that of the
single star previously studied by
García-Berro et al. (1999), the only difference being that the
second dredge-up is not caused by carbon burning, but, instead, by the
gravothermal energy release at the end of core helium burning and before
carbon is ignited off-center in the degenerate core.
Another important difference between case 1.a and case 1.b is the evolutionary stage of the secondary at the time at which reversal mass transfer occurs. For the case 1.a, it occurs when the secondary component climbs the red giant branch, whereas for case 1.b it only fills its Roche lobe when it reaches the Super-AGB phase. However, the uncertainties related to stellar winds, both in mass-loss rates and in the efficiency of mass accretion, do not allow us to determine the exact orbital parameters at the begining of reversal mass transfer and, ultimately, whether merger episodes could occur. In fact, if mass loss from the secondary (as a red giant or as a Super-AGB star) allows a significant decrease in its mass with low enough mass transfer rates, the steep change in the density profile of the mass losing star at the border of its core could brake the spiral-in of the primary. The system would then be able to survive as a binary, and the following possibilities may arise:
If, on the other hand, accretion onto the primary is efficient
enough so that this component is able to reach the Chandrasekhar
mass, it will become a neutron star after a supernova explosion
(Gutiérrez et al. 1996). Furhermore, if during the accretion
phase of the primary, the secondary does not lose a significant
amount of mass, our system may help to explain the HMXB precursor
PSR J1740-3052 (Stairs et al. 2001), a binary system with
days, composed of a neutron star plus a massive (
)
companion. Similarly, the highly eccentric
pulsar PSR B1259-63 (Johnston et al. 1992) might also be another
possible observational counterpart, as the non-degenerate
companion of the neutron star has a mass
.
Further evolution of the secondary might also lead it to undergo a supernova explosion and leave a second neutron star as a remnant, and hence a binary pulsar could also be a possible outcome. Such a system may remain bound or be disrupted, depending on the asymmetry of the supernova explosion and on the amount of mass ejected from the system. A possible observational counterpart could be the binary pulsar PSR B1813+16 (Taylor et al. 1976).
Acknowledgements
Part of this work was supported by the Spanish DGES project number PB98-1183-C03-02, by the MCYT grant AYA2000-1785, and by the CIRIT. We also wish to thank J. José for carefully reading the manuscript, and to D. Vanbeveren, for his numerous suggestions.