A&A 371, 152-173 (2001)
DOI: 10.1051/0004-6361:20010309
P. Marigo1 - L. Girardi1 - C. Chiosi1 - P. R. Wood2
1 - Dipartimento di Astronomia, Università di Padova,
Vicolo dell'Osservatorio 2, 35122 Padova, Italia
2 -
Mount Stromlo and Siding Spring Observatories, Australian National
University, Private Bag, Weston Creek PO, ACT 2611, Australia
Received 12 December 2000 / Accepted 28 February 2001
Abstract
We present extensive evolutionary models of stars
with initial zero-metallicity, covering a large range of initial masses
(i.e.
). Calculations are carried
out at constant mass, with updated input physics, and applying an
overshooting scheme to convective boundaries. The nuclear network
includes all the important reactions of the p-p chain,
CNO-cycle and
-captures, and is solved by means
of a suitable semi-implicit method. The evolution is followed up to the
thermally pulsing AGB in the case of low- and intermediate-mass stars,
or to the onset of carbon burning in massive stars.
The main evolutionary features of these models are discussed,
also in comparison with models of non-zero metallicity. Among
several interesting aspects, particular
attention has been paid to describe: i) the first synthesis of
12C inside the stars, that may suddenly trigger the CNO-cycle
causing particular evolutionary features; ii) the pollution of
the stellar surface by the dredge-up events, that are effective
only within particular mass ranges; iii) the mass limits which conventionally
define the classes of low-, intermediate-, and high-mass stars on the
basis of common evolutionary properties,
including the upper mass limit for the achievement
of super-Eddington luminosities before C-ignition in the high-mass regime;
and iv) the expected pulsational properties of zero-metallicity stars.
All relevant information referring to the evolutionary tracks
and isochrones is made available in computer-readable
format.
Key words: stars: evolution - stars: interiors - stars: Hertzsprung-Russell (HR) diagram - stars: low-mass
The stars with negligible, if not zero, initial metal abundance - named population III objects (Pop-III for short) - are expected to have structure and evolution distinct from those in which traces of heavier elements, like those of the CNO group, are present.
The major difference between zero-metal stars and those of
normal metal content (even down to very low values of Z) lies in the
mechanism of nuclear energy generation. In fact, owing to the lack of
CNO nuclei, the pre-main sequence gravitational contraction cannot
stop until the central temperature and density
are high enough to allow the p-p chain to provide the energy budget.
Since the p-p chain is a poor thermostat as compared to the CNO-cycle,
very high temperatures can be reached in the central regions. This
may eventually lead (depending on the mass of the star)
to the first synthesis of primary carbon through the 3-
reaction, while the star is still on the main sequence.
As a consequence, the CNO-cycle is activated
at the very high temperatures characterising the 3-
process,
possibly causing a dramatic change
in the dominant energy source (from p-p chain to CNO-cycle), hence
affecting the stellar structure amid the main sequence evolution.
As demonstrated below, similar changes in the
evolutionary behaviour may also occur at later stages, whenever carbon
is first produced (or transported) into different regions of the star.
The threshold abundance of
at which the distinct behaviour
of zero metal stars appears is as low as
(see Cassisi & Castellani 1993).
Starting from these values and higher, the efficiency of
CNO-cycle is sufficient to recover the standard behaviour.
Needless to say, the distinct evolutionary behaviour of these stars may also imply a very distinct nucleosynthesis and chemical pollution of the interstellar medium, if compared to stars with initial Z>10-10. Also their properties in the HR-diagram should be quite peculiar. Both aspects are particularly important in the modelling of
The plan of the paper is as follows.
After briefly reviewing the available literature
on evolutionary models of primordial stars (Sect. 2),
the main assumptions of present calculations are described in
Sect. 3 concerning the nucleosynthesis,
and Sect. 4 for the other physics prescriptions.
Then, the presentation and discussion of the results start with
Sect. 5 - dealing with general evolutionary aspects
(e.g. the onset of the 3-
reaction and the activation
of the CNO cycle, lifetimes, etc.) - and proceed through
Sects. 6 and 7, which focus on particular
issues of low- and intermediate-/high-mass models,
respectively (e.g. the AGB phase, the Eddington critical luminosity).
Our predictions for possible changes in the surface chemical
composition are given in Sect. 8.
Finally, Sect. 9 is dedicated to an analysis of
the expected pulsational
properties of Z=0 stars with different masses and evolutionary stages,
to investigate whether pulsation might be a discriminating
feature in observational searches of Pop-III objects.
Appendices describe the tables of the Z=0 evolutionary tracks
and related isochrones available in electronic format.
However, over the years,
interest has turned towards more "normal'' stars.
In fact, various fragmentation models of primordial gas have suggested that
the first stars might have formed with low/intermediate masses, e.g.
Yoshi & Saio (1986) find the peak of the mass function at roughly
;
Nakamura & Umemura (1999) find a the typical
mass of
,
which may grow to
by accretion;
and Nakamura & Umemura (2000) find a lower mass-limit extending
down to
.
With more detailed and extended
calculations for the collapse and fragmentation of primordial gas
clouds, Nakamura & Umemura (2001) suggest that the
initial mass function of Pop-III stars may
be bimodal with peaks at about
and
,
the relative heights of these peaks being a
function of the collapse epoch.
The provisional conclusion we may get from all this is that
the mass peak of Pop-III stars is likely greater than say
and
possibly in the range
1 to
.
However, both lower and much higher values
cannot be firmly excluded.
Ezer & Cameron (1971) first studied models of Pop-III stars in
the mass range 5 to
from the
pre-main sequence until core H-exhaustion.
In a subsequent paper Ezer (1972) investigated the evolution of a
star up to the stage of core He-exhaustion.
Cary (1974) and Castellani & Paolicchi (1975) computed
zero-age-main-sequence (ZAMS)
models in the mass range 2 to
and 1 to
,
respectively.
Woosley & Weaver (1981) followed the evolution at constant mass
of a
star until the supernova explosion stage, thus
allowing the first estimate of the contamination power of this type of star.
Very massive Pop-III stars were first investigated by Bond et al.
(1982, 1983).
El Eid et al. (1983) and Ober et al. (1983) explored the mass range
from 80 to
,
including
mass loss by stellar winds up to the stage of supernova explosion or
collapse to a black hole. Similar studies were undertaken by Klapp (1983,
1984) for very massive stars in the range 500 to
.
These papers
made it possible to get an estimate of the chemical yields both in the
wind and explosive stages.
Forieri (1982) studied the evolution of massive
Pop-III stars in mass range 10 to
up to the central
He-exhaustion stage. These models were evolved at constant mass
but included convective overshoot.
Castellani et al. (1983) presented models for 10, 15 and
stars from ZAMS to C-ignition.
The advanced phases of low-mass stars remained almost unexplored
till the early study
by D'Antona (1982) of a
star up to the stage of core
He-flash, where for the first time the much less extended RGB was noticed
(see also Castellani & Paolicchi 1975 for an earlier suggestion).
Chieffi & Tornambé (1984) investigated the evolution of a
object, noticing the absence of thermal instabilities
during the double-shell (AGB) phase.
Fujimoto et al. (1984) pointed out that threshold values for
the core mass and CNO abundances exist for the occurrence of
He-shell flashes during the AGB phase.
Cassisi & Castellani (1993) presented grids of models with masses in the
range
and metallicity from
Z=10-4 down to
Z=10-10. The models were followed
from ZAMS till He-flash, final cooling of the CO core
or C-ignition as appropriate. Cassisi et al. (1996) analysed
the RR Lyrae properties and the unusually strong He-shell flashes
of low-mass stars with very low initial metal content
(
).
Many recent/ongoing work on zero-metallicity stars have been recently presented in the symposium proceedings edited by Weiss et al. (2000a), to which the reader should refer for a comprehensive picture of the present state of understanding of these objects. Particularly impressive are the studies being carried out by Heger et al. (2000) and Umeda et al. (2000), where massive zero-metallicity stars are made to evolve up to the pre-supernova stage and supernova explosion, respectively. These studies describe in detail the nucleosynthesis of heavy elements occurring in these stars.
In the work by Weiss et al. (2000b) the authors investigate the role of
diffusion and external pollution in the evolution of metal-free
low-mass stars, finding that no sizable effects occur (the accreted
metals cannot reach the central H-burning regions,
which would have otherwise affected the structure of zero-metallicity stars).
Using an extended nuclear network, they also demonstrate
that hot p-p chains are less effective than the 3-
process in
producing carbon inside a zero-metallicity star.
Finally, we should mention the work by Schlattl et al. (2001), who investigate the possible changes in the envelope chemical composition as the result of convective dredge-up during the He-flash of Pop-III RGB stars (see also Fujimoto et al. 2000).
In this work the adopted nuclear network consists of the
sets of reactions for the pp1-chain
The simultaneous consideration
of H- and He-burning reactions is required by the fact that
in these stars H-burning may occur at temperatures and densities
high enough to allow the synthesis, via -captures, of the
missing catalysts (i.e. 12C and 16O)
for the activation of the CNO cycle.
The choice of a semi-implicit method is motivated by the search for a convenient compromise between the higher accuracy typical of the explicit scheme, and the better stability of the solution guaranteed by the implicit scheme. We recall that, from a technical point of view, in implicit/explicit methods the increments of the dependent variables over a given integration step are calculated from the derivatives evaluated at the new/old location of the independent variable. The net effect is that implicit methods are stable for large integration steps, but with a certain loss of accuracy in following the solution towards the equilibrium. The reverse situation occurs for explicit methods.
Actually, implicit methods are required when dealing with stiff differential equations, in which the changes of the dependent variables involve quite different scales of the independent variable. This feature matches the case of the abundance equations, as different elemental species are characterised by very different nuclear lifetimes. For instance, at temperatures near 25 106 K, the e-folding time for 15N is on the order of years, whereas for 14N it is on the order of 105 years. In general, implicit methods converge to the true equilibrium solutions of the nuclear network for relatively large time steps, whereas purely explicit methods would require extremely short time steps to both find the solution and maintain stability. Alternatively, in order to make the explicit scheme less time consuming this may be coupled with external assumptions on nuclear equilibria (usually made for the CNO nuclei). This possible choice is however quite risky, specially in the case of zero-metallicity stars in which the total abundance by number of the CNO elements during H-burning is not constant, but rather increasing thanks to the contribution from He-burning.
The energy losses by pair, plasma, and bremsstrahlung neutrinos, are from Haft et al. (1994).
The extension of convective boundaries is estimated by means of an
algorithm which takes into account overshooting from
both core and envelope convective zones. The formalism is fully
described in Bressan et al. (1981) and Alongi et al. (1991).
The main parameter describing
overshooting is its extent
across
the border of the convective zone, expressed in units of pressure
scale height.
We adopt the following prescription for the
parameter
as a function of stellar mass:
Overshooting at the lower boundary of
convective envelopes is also considered.
The value of
(see Alongi et al. 1991, for a description of the formalism) is adopted for
stars with
.
For
M>2.0
a value of
is assumed.
Perhaps the most uncertain aspect of the evolution of Pop-III stars is mass loss by stellar winds. Two typical situations can be envisaged: (i) Mass loss from massive stars during their whole evolutionary history, and (ii) mass loss during the RGB and AGB phases of low and intermediate-mass stars.
Of course, the first question is whether our knowledge for normal-metallicity stars can be extrapolated and safely used for Pop-III stars. This depends on the physical nature of the mechanism driving mass-loss. In massive stars of normal metallicity (no matter whether low or high) mass loss at high effective temperature is due to radiation pressure on resonance lines of ions like CIV, NV, SiIV, etc., whereas that at low temperatures is likely caused by radiation pressure on dust grains and/or due to pulsational properties (see Chiosi & Maeder 1986; Chiosi et al. 1992; Maeder & Conti 1994; Feast 1991; Kudritzki 1998; for recent reviews). Furthermore, the body of observational and theoretical information allows us to establish empirical relations according to which the mass loss rates are seen to decrease with the metallicity. On this basis one could perhaps argue that as long as the surface layers keep the original composition mass loss cannot occur. If and when the surface abundances are contaminated by mixing of heavy elements, mass loss by the above mechanisms could be activated. However, the above reasoning would immediately fail if other mechanisms (metallicity independent) concur to drive mass loss.
Even for non-zero metallicities, modelling the RGB and AGB phases of low- and intermediate-mass stars is usually problematic, as we do not have a definitive theoretical scenario to account for mass loss during these stages. Moreover, the question arises as to whether the mass-loss prescriptions normally used for metal-poor globular cluster giants (e.g. Reimers 1975), or for pulsating AGB stars (e.g. Vassiliadis & Wood 1993) can be applied to Pop-III stars. Anyway, if any of these empirical relations is actually used in zero-metallicity models, very little mass is found to be lost during the RGB and AGB phases. This justifies to some extent the assumption of evolution at constant mass, that we adopt in the present paper.
However, in Sect. 7 we explore the possibility that
the more massive and evolved Pop-III stars may become
gravitationally unbound and hence lose mass. These stars will be
the subject of a forthcoming paper (Marigo et al. 2001, in preparation).
![]() |
Figure 1:
Zero-metallicity stellar
tracks in the HR diagram, for initial masses in the range
0.7 ![]() ![]() ![]() |
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For the sake of clarity in the presentation of the
results, we adopt the standard convention of sub-grouping the models
as a function of the initial mass,
by defining two mass limits, namely
and
.
For low-mass stars, the evolution is interrupted at the stage of
He-flash in the electron-degenerate hydrogen-exhausted
core. The evolution is then re-started from a
Zero Age Horizontal Branch (ZAHB) model having the
same core mass
and envelope
chemical composition as the last RGB model (see also Sect. 6.3).
Nuclear lifetimes of the H- and He-burning phases as a function
of the stellar mass are reported in Table 1.
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0.7 | 2.2751 1010 | 1.1598 108 |
0.8 | 1.3749 1010 | 1.1536 108 |
0.9 | 8.9023 109 | 1.0576 108 |
1.0 | 6.1077 109 | 1.1029 108 |
1.1 | 4.3791 109 | 2.0626 108 |
1.2 | 3.3170 109 | 1.5657 108 |
1.3 | 2.5358 109 | 1.2843 108 |
1.4 | 1.9741 109 | 1.0660 108 |
1.5 | 1.5810 109 | 9.4372 107 |
1.6 | 1.2830 109 | 7.5237 107 |
1.7 | 1.0510 109 | 6.3585 107 |
1.8 | 8.8129 108 | 6.2271 107 |
1.9 | 7.4387 108 | 4.8001 107 |
2.0 | 6.3454 108 | 4.2688 107 |
2.1 | 5.4827 108 | 3.7334 107 |
2.2 | 4.7794 108 | 3.5347 107 |
2.3 | 4.2032 108 | 3.0438 107 |
2.4 | 3.7470 108 | 2.6636 107 |
2.5 | 3.3439 108 | 2.4283 107 |
2.7 | 2.7806 108 | 2.0256 107 |
3.0 | 2.0931 108 | 1.6214 107 |
3.5 | 1.4855 108 | 1.2113 107 |
4.0 | 1.0755 108 | 1.1333 107 |
5.0 | 6.7291 107 | 6.9363 106 |
6.0 | 4.7335 107 | 4.4815 106 |
6.5 | 4.1719 107 | 3.5485 106 |
7.0 | 3.5567 107 | 2.9240 106 |
8.0 | 2.6947 107 | 2.0781 106 |
8.3 | 2.5305 107 | 1.9528 106 |
9.0 | 2.2221 107 | 1.6740 106 |
9.5 | 2.0475 107 | 1.5268 106 |
10.0 | 1.8923 107 | 1.4089 106 |
12.0 | 1.4710 107 | 1.0536 106 |
15.0 | 1.1307 107 | 7.7869 105 |
20.0 | 8.4720 106 | 5.7555 105 |
30.0 | 5.9030 106 | 4.2264 105 |
50.0 | 4.1315 106 | 3.3023 105 |
70.0 | 3.4538 106 | 2.9938 105 |
100.0 | 2.9509 106 | 2.6841 105 |
Let us now discuss the main features of the models here presented. It should be recalled that many aspects of the evolution of zero-metallicity stars are already described by Cassisi & Castellani (1993; hereinafter also CC93). In order to avoid redundant discussion, the present analysis is more focused on particular issues and/or open questions which, to our opinion, still deserve further investigation.
First, we can see that, as expected, the Z=0 tracks
in the
diagram are
systematically hotter and more luminous.
The latter feature results in shorter nuclear lifetimes
(given in Table 1) with respect to those of solar-composition
models.
Other obvious differences refer to (i) some morphological features of the tracks, and (ii) the relative excursion in in effective temperature characterising the nuclear H- and He-burning phases.
Concerning point (i), we notice that well-defined
loops may show up on the Z=0 tracks (e.g. 1
and
3
models in Fig. 2).
The physical interpretation of these
unusual features will be discussed in Sect. 6.
Concerning point (ii), it is worth making the following remarks.
As we can see, the locus of points
corresponding to the termination of the main sequence (squares)
describes a path in the
plane, which significantly
bends towards lower effective temperatures at increasing stellar mass
(i.e.
).
This can be explained as due to the gradual
decrease of the mass exponent a in the mass-luminosity relation
(i.e.
), as pointed out long ago by Stothers (1966).
Comparing the evolution of
models with Z=0 and Z=0.0004,
we see that once they experience powerful He-ignition at the tip
of the RGB, the subsequent quiescent core He-burning occurs
either close to the Hayashi line (Z=0.0004), or
to an HB structure blueward of the RR Lyrae instability strip
(extending from
to
;
see Bono et al. 1995) when Z=0.
Moreover, in stars with
,
the onset of central He-burning (circles)
occurs soon after hydrogen exhaustion
in the core, and proceeds as these stars are evolving
towards lower effective temperatures at nearly constant
luminosities. Into this regard, a feature of particular interest
is that the Z=0 models of intermediate mass
(
)
start and complete their
core He-burning phase always in the bluest parts of the
tracks, i.e. quite far from their
Hayashi lines. This implies, for instance, that Z=0 intermediate-mass
stars would never appear as either red-clump stars or blue-loop stars.
For even more massive models,
the locus of points marking the end of the He-burning
phase (starred symbols in Fig. 2) also presents a systematic
bending towards cooler regions at increasing stellar masses.
In particular, we find that massive models with
may reach their
Hayashi lines already during the helium-burning phase,
remaining there until central carbon ignition.
In general, it should be remarked that whether or not a star spends some part of its evolution on (or in proximity to) the Hayashi line may have important implications. In fact, the corresponding development of a deeply extended convective envelope may cause dredge-up events, with consequent changes in the surface chemical composition. This point will be discussed in Sect. 8.
In addition to the
plane,
another powerful diagnostic of the stellar structure is
the central temperature - central density diagram, as
it illustrates the state of the gas in the innermost regions.
Figure 3 displays the evolutionary behaviour
in this plane by the models here considered.
Lines are drawn to mark the approximate boundaries of the regions in which
the equation of state is dominated by perfect gas, degenerate electron gas,
and radiation pressure.
Point loci as a function of the stellar
mass are shown for relevant evolutionary stages.
One can notice, in this figure, some typical behaviours
that depend mainly on the equation of state, and that are found in
stellar models of any metallicity. For instance,
in correspondence to the stages of gravitational contraction of
stellar cores between nuclear burnings,
the evolution of the central conditions in the domain of ideal gas
can be described by
,
as expected from homology considerations (e.g. Kippenhahn & Weigert
1990).
However, the
diagram of zero-metallicity stars
contain several unusual features worthy of note, such as:
the peculiar loci marking the
onset of central H-burning (line a) and the onset of the 3-
reaction
and the CNO-cycle (line b),
as well as some short-lived periods of core expansion and/or heating
(that are evident, for instance, soon above line b), and the unusually
low mass limit (
)
for stars to become degenerate before
He-ignition. Each one of these features will be discussed below.
![]() |
Figure 2:
Zero-metal evolutionary tracks (solid lines)
for selected initial masses (in ![]() ![]() |
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Figure 3:
Evolution of central conditions (i.e. density
and temperature) for the whole set
of Z=0 models. Approximate relations are used to determine
the boundaries (thick straight lines)
of the regions in which the equation of state is expected to be
dominated by different pressure components, i.e. degenerate gas, perfect
gas, and radiation pressure.
Labelled lines indicate point loci as a function
of the stellar mass corresponding to the onset of central H-burning (a);
onset of the 3-![]() ![]() ![]() |
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The onset of the p-p reactions as a function of mass is
very inclined in the
diagram (line a in
Fig. 3).
Due to the well known rather weak dependence
of the p-p reaction rates on
temperature (i.e.
),
the contraction of the central regions can easily proceed,
so that high temperatures may be then attained during the H-burning
phase.
Once the central temperatures have increased up to typical values
of
,
the ignition of the 3-
reaction takes place.
This occurrence marks a fundamental event, as it is represents the
first significant production of metals in these stars.
As a consequence, the synthesis of primary
12C leads to the activation of the CNO-cycle, which
then starts providing nuclear energy in competition with
the p-p reactions (line b in Fig. 3).
The onset of the 3-
reaction occurs at earlier
and earlier stages at increasing stellar mass (see Fig. 4).
This occurs essentially because the 3-
reaction
requires much higher temperatures than the p-p chain.
In the lowest mass models, the first
production of 12C takes place towards the very end of the
H-burning phase, and may not even occur for
.
In this latter case, all central hydrogen is entirely burnt via the
p-p chain.
At a sufficiently high mass (
in our models),
the 3-
reaction ignites even before
the p-p reactions have slowed down the initial stellar contraction.
In these models then, H-burning simply proceeds via the CNO-cycle,
without any significant phase of central burning via the p-p chain.
The above features explain the striking change in the slope of the
curve corresponding to the onset of the H-burning phase as a
function of the stellar mass (line a in Fig. 3).
For models with
,
H-burning
starts above a nearly vertical line in the
diagram (i.e. at varying
and for an almost constant
), that is characteristic of ZAMS
stars that burn hydrogen predominantly through the p-p chain.
For models with
,
H-burning starts
above a line of modest slope (i.e. a modest increase
of
corresponds to a substantial decrease of
),
that simply traces the minimum
and
necessary to ignite the 3-
reaction (and hence the CNO-cycle) in a core with
.
This also explains why line (a) simply merges with line (b) at
.
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Figure 4: The central abundance of hydrogen when the fractional abundance of newly synthesised CNO elements grows up to 10-14 |
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Subsequent to the first carbon production,
as the energy contribution of the CNO reactions increases
(at the expense of the p-p reactions) the central regions are forced to
expand, with a consequent decrease of the central density (and
central temperature in the most massive models). This
leads to the behaviour of the tracks in Fig. 3,
around the line (b).
The activation of the CNO-cycle
may result in the appearance of some peculiar features in the H-R diagram
of low-mass models (i.e. loops; see Sect. 6) and,
given its quite large temperature dependence
(i.e.
),
it also favours the development
or a larger extension of core convection.
For instance, in the case of
the models with
displayed in
Fig. 5, after the initial recession
of the convective core, central H-burning via the p-p
reactions occurs in radiative conditions.
Then, as the CNO cycle turns on,
central regions become convectively unstable again.
This effect is anticipated at increasing stellar mass, and eventually
in more massive models (
)
the transition
of the core from convective to radiative and back to convective
does not occur at all.
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Figure 5:
Evolution of core convection during central H-burning
for a few selected models with masses (in ![]() |
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From the location of the non-degenerate/degenerate boundary
and of the curves for He- and C-ignition shown
in Fig. 3,
it is possible to estimate the critical masses
and
.
Figure 6 illustrates the predicted trend of the
mass limits
and
as a function of the metallicity,
combining the results by Girardi et al. (2000) with those of the present
work for Z=0.
Concerning
,
this is found to decrease with the metallicity,
dropping down to
for Z=0.
This is the combined result of the larger convective cores during
the MS phase at decreasing metallicities
when the CNO cycle is the dominant energy source
(up to
,
see CC93)
and, at even lower Z, of the higher temperatures
reached in the centre when H-shell burning mainly occurs via the p-p
reactions (as is the case for the Z=0 models).
Concerning
,
its trend is not monotonic as it first
decreases with metallicity, reaches a minimum, and finally increases
again up to
when Z=0.
This behaviour is essentially controlled by the mass of the He-core
left at the end of the MS.
For CNO-dominated H-burning, lower metallicities
correspond to more concentrated energy sources and
larger convective cores, which explains the initial
decrease of
.
As soon as the energy contribution from the p-p reactions becomes
competitive with the CNO cycle, the central burning regions on the MS
are more extended, while the sizes of the convective cores
tend to be smaller. This explains the presence of the minimum
and the subsequent increase of
at decreasing Z.
![]() |
Figure 6:
Critical masses
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Both behaviours of
and
are in agreement with
the results of CC93. Their larger values of
and
at given metallicity essentially reflect
the different treatment of convective boundaries (i.e. Schwarzschild
criterion) with respect to ours (i.e. overshoot scheme;
see Sect. 4).
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Figure 7:
Kippenhahn diagrams showing the location (in fractional mass)
of the burning and convective zones during the evolution of the
![]() |
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Another remarkable evolutionary feature is
found in stars with masses
-2.5
,
and
.
During the central He-burning phase, the outer boundary of the
core overshooting zone is located very close to the bottom of
the H-burning shell.
Initially, we found that, if the convective core was left to grow, it eventually reaches the H-shell and engulfs some H-rich material, which is rapidly burnt in the core via the CNO cycle. This causes a flash that expands the core, so that central He-burning weakens and the convective core recedes temporarily (in mass). After the flash has occurred, the convective core starts growing again, possibly approaching the bottom of the H-shell. This leads to quasi-periodic H-flashes, the related growth/recession of the core resembling some sort of breathing convection.
However, this picture is likely to be physically unsound because of the treatment adopted in the evolution code. In fact, the H-flash is caused by the fact that the engulfed hydrogen is first mixed throughout the convective core, and then burnt according to nuclear reaction rates which are mass-averaged all over the convective layers. This procedure is a good approximation as long as the convective lifetimes are shorter than the nuclear lifetimes, which is usually the case for CNO-burning occurring in "ordinary'' conditions.
But at the very high temperatures
(108 K) reached at the bottom the H-burning shell,
the nuclear lifetimes involved in the CNO cycle (on the order of hours, days)
can be comparable or even shorter than the typical
convective timescales (on the order of months).
Under these conditions, a correct approach would require a
time-dependent solution scheme which couples simultaneously
nucleosynthesis and mixing (see e.g. Schlattl et al. 2001).
Leaving the full analysis of this point to a future investigation,
we make the reasonable assumption in the present study
that hydrogen burns locally at the bottom
of the H-shell, before a complete mixing by core convection can occur.
Technically, this translates into the
condition that the maximum allowed extension of core convection
is set by the bottom of the H-burning shell.
Figure 7 illustrates the evolution of the convective and
burning regions of the 2.2
model.
We can notice the approach of core convection
towards the inner location of the
H-burning shell, during the He-burning phase (third panel from left).
We can also notice the remarkable thickness (in mass) of the
H-burning shell, which extends deeply inward during both shell
H-burning and core He-burning phases.
This is due to the fact that hydrogen in the shell
is mainly burnt via the p-p chain, which due to its low
temperature sensitivity, allows quite high temperatures to be
attained at the bottom of the shell.
At an age of 5.1 108 yr, our numerical treatment starts to limit the extension
of the core convection. In this way, the evolution can
be quietly followed up to the end of He-burning.
![]() |
Figure 8:
Physical properties of the 1 ![]() ![]() ![]() ![]() ![]() |
Open with DEXTER |
As soon as a CNO abundance of about
10-10,-9 is built at the centre
of the star, a sudden spike arises in the rate of energy generation
by nuclear H-burning. This occurs when the growing efficiency of the CNO
cycle exceeds that of the p-p reactions (
),
which have dominated the energetics of the star up to then.
The central conditions react to this sudden increase of energy production:
the innermost regions expand so that both
and
decrease.
![]() |
Figure 9: H-R tracks of Z=0 low-mass models extending from the ZAHB up to the beginning of the TP-AGB phase |
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M |
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|
0.7 | 0.5015 | 2.573 | -0.0325 |
0.8 | 0.4991 | 2.549 | -0.0322 |
0.9 | 0.4951 | 2.497 | -0.0315 |
1.0 | 0.4882 | 2.400 | -0.0303 |
1.1 | 0.3683 | 0.989 | -0.0100 |
We have computed this
quantity for each evolutionary track, by simply assuming
that the flash provides just the energy
necessary to lift the electron degeneracy from the core.
In practice, this energy
can be evaluated as the difference in total energy between the last
RGB-tip configuration and the initial ZAHB one,
.
Then,
the mass of burnt helium, and hence
,
follows straightforwardly from the Q-value of 3-
reactions.
Table 2 presents some relevant
characteristics of the models at the stages of He-flash and ZAHB.
We can notice that the predicted fractional abundance of helium
is burnt in the flash amounts to few percent (typically 3%).
The subsequent evolution on the HB (see Fig. 9)
is quite similar to that of
Pop-II models. However, as already pointed out by Cassisi et al. (1996),
all low-mass Z=0 stars burn their central He at
quite high effective temperatures (with
),
so that a negligible fraction of RR Lyrae pulsators on the ZAHB
is expected.
This aspect is further commented on Sect. 9.
The evolutionary properties of metal deficient AGB stars have been analysed in the past by Fujimoto et al. (1984); Chieffi & Tornambé (1984); Cassisi et al. (1996), and more recently by Domínguez et al. (1999). A crucial question addressed in these works is about the possible occurrence of the He-shell flashes (thermal pulses), which are usually found in AGB models of various (but not zero) metallicities.
According to the semi-analytical investigation of Fujimoto et al. (1984)
on the stability of nuclear burning shells,
thermal pulses are conditional on the existence of thresholds in
core mass ,
and CNO abundance
in the envelope.
To summarise, He-shell flashes are expected to occur when
i)
for any value of
,
and ii)
if
,
the critical lower limit
depending on
.
The dependence on
can be qualitatively understood considering
that in the absence (or strong deficit) of CNO elements in the envelope,
the H-burning shell itself
produces new carbon via the 3-
reaction,
so that the CNO cycle can operate. It follows that the H-
and He-burning shells are not energetically decoupled, and proceed outward
in mass at the same rate (see also Chieffi & Tornambé 1984).
Since He-rich material is not accumulated in the inert intershell buffer,
but rather steadily burnt by the underlying He-shell,
the occurrence of thermal pulses is thus prevented.
In this study we present the results for two intermediate-mass models,
with initial zero metallicity and
masses of
and
,
which are made evolve through the initial stages of the AGB phase,
including a few episodes of thermal instabilities of the He-burning
shell. Whether such events may be ascribed to true He-shell flashes
is discussed below.
As the
and
models reach their Hayashi tracks (E-AGB phase)
after central helium exhaustion, envelope convection progressively
moves inward in both cases, and extends into regions which have
experienced earlier nucleosynthesis.
The so-called "second dredge-up'' takes place, but with a substantial
difference between the two models. This can be appreciated by looking at
Table 3, which presents the predicted surface chemical
compositions after the second dredge-up event for all our models.
In the
model the envelope penetrates into the
chemical profile left by the recession of the convective core during
the MS, and into regions previously affected by the p-p chain,
so that the surface enrichment is essentially in helium
with no trace of heavier metals
(see also Fig. 7 for a similar case).
In the
model the envelope penetration
proceeds further, reaching down
to the former location of the H-burning shell
(which temporarily extinguishes), so that
newly synthesized CNO elements are also
dredged-up to the surface.
After the occurrence of the second dredge-up, the evolution of
these AGB models proceeds quite differently.
![]() |
H | 3He | 4He | 12C | 13C | 14N | 15N | 16O | 17O | 18O | 22Ne | 24+25Mg | 20Ne |
Initial: | |||||||||||||
all | 0.770 | 2.46 10-5 | 0.230 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
After the first dredge-up: | |||||||||||||
0.7 | 0.769 | 8.63 10-5 | 0.231 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
0.8 | 0.768 | 1.40 10-4 | 0.232 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
0.9 | 0.768 | 1.54 10-4 | 0.232 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
1.0 | 0.768 | 1.17 10-4 | 0.232 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
1.1 | 0.769 | 8.29 10-5 | 0.231 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
After the second dredge-up: | |||||||||||||
1.2 | 0.769 | 6.49 10-4 | 0.231 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
1.3 | 0.767 | 4.89 10-4 | 0.233 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
1.4 | 0.769 | 4.72 10-4 | 0.231 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
1.5 | 0.767 | 3.45 10-4 | 0.233 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
1.6 | 0.767 | 3.16 10-4 | 0.232 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
1.7 | 0.767 | 2.75 10-4 | 0.233 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
1.8 | 0.769 | 2.64 10-4 | 0.231 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
1.9 | 0.768 | 2.19 10-4 | 0.232 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
2.0 | 0.768 | 2.02 10-4 | 0.232 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
2.1 | 0.704 | 1.07 10-4 | 0.296 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
2.2 | 0.697 | 9.59 10-5 | 0.303 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
2.3 | 0.691 | 8.13 10-5 | 0.309 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
2.4 | 0.690 | 6.82 10-5 | 0.310 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
2.5 | 0.681 | 5.97 10-5 | 0.319 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
2.7 | 0.669 | 5.04 10-5 | 0.331 | 4.17 10-21 | 1.34 10-21 | 4.95 10-19 | 2.02 10-23 | 1.28 10-20 | 1.33 10-22 | 7.16 10-26 | 2.43 10-35 | 1.62 10-50 | 7.70 10-38 |
3.0 | 0.660 | 4.10 10-5 | 0.340 | 1.13 10-18 | 3.62 10-19 | 1.28 10-16 | 5.19 10-21 | 3.15 10-18 | 2.82 10-20 | 1.60 10-23 | 4.25 10-32 | 6.91 10-46 | 3.19 10-34 |
3.5 | 0.649 | 2.99 10-5 | 0.351 | 8.06 10-16 | 2.59 10-16 | 8.58 10-14 | 3.46 10-18 | 1.97 10-15 | 1.41 10-17 | 8.69 10-21 | 2.06 10-28 | 1.86 10-40 | 6.19 10-30 |
4.0 | 0.644 | 2.29 10-5 | 0.356 | 6.33 10-14 | 2.03 10-14 | 6.17 10-12 | 2.46 10-16 | 1.31 10-13 | 7.41 10-16 | 5.23 10-19 | 5.33 10-26 | 7.42 10-37 | 4.13 10-27 |
5.0 | 0.637 | 1.56 10-5 | 0.363 | 3.58 10-9 | 7.81 10-12 | 6.06 10-10 | 1.57 10-14 | 4.54 10-12 | 3.90 10-14 | 5.11 10-15 | 5.22 10-19 | 6.68 10-28 | 1.14 10-23 |
6.0 | 0.632 | 1.11 10-5 | 0.368 | 2.63 10-7 | 1.19 10-11 | 9.35 10-10 | 2.41 10-14 | 7.36 10-11 | 5.91 10-14 | 4.67 10-13 | 2.96 10-16 | 1.66 10-23 | 4.73 10-21 |
6.5 | 0.626 | 9.25 10-6 | 0.374 | 7.07 10-7 | 1.46 10-11 | 1.13 10-9 | 2.90 10-14 | 2.99 10-10 | 7.37 10-14 | 2.21 10-12 | 1.95 10-15 | 2.63 10-22 | 4.09 10-20 |
7.0 | 0.668 | 8.74 10-6 | 0.332 | 1.26 10-13 | 4.04 10-14 | 1.09 10-11 | 4.30 10-16 | 2.08 10-13 | 8.57 10-16 | 7.51 10-19 | 6.02 10-26 | 2.75 10-36 | 9.03 10-27 |
8.0 | 0.681 | 7.49 10-6 | 0.319 | 9.66 10-14 | 5.68 10-17 | 8.08 10-12 | 6.32 10-19 | 1.47 10-13 | 1.41 10-18 | 1.15 10-21 | 9.09 10-30 | 6.05 10-42 | 4.43 10-31 |
8.3 | 0.687 | 7.00 10-6 | 0.313 | 1.19 10-14 | 3.81 10-15 | 1.01 10-12 | 3.98 10-17 | 1.88 10-14 | 7.39 10-17 | 6.71 10-20 | 1.95 10-27 | 2.12 10-38 | 2.36 10-28 |
9.0 | 0.770 | 4.30 10-5 | 0.230 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
9.5 | 0.770 | 4.30 10-5 | 0.230 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
10.0 | 0.770 | 4.25 10-5 | 0.230 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
12.0 | 0.770 | 3.63 10-5 | 0.230 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
15.0 | 0.770 | 2.47 10-5 | 0.230 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
20.0 | 0.770 | 1.66 10-5 | 0.230 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
30.0 | 0.770 | 1.02 10-5 | 0.230 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
50.0 | 0.770 | 5.95 10-6 | 0.230 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 | 0 |
70.0 | 0.713 | 1.06 10-6 | 0.287 | 7.65 10-12 | 2.45 10-12 | 6.39 10-10 | 2.53 10-14 | 1.30 10-11 | 5.29 10-14 | 5.34 10-17 | 5.11 10-24 | 2.08 10-14 | 2.58 10-14 |
100.0 | 0.640 | 6.39 10-7 | 0.360 | 2.27 10-11 | 7.25 10-12 | 1.87 10-9 | 7.21 10-14 | 3.53 10-11 | 1.33 10-13 | 1.33 10-16 | 2.03 10-23 | 1.42 10-32 | 1.13 10-23 |
As we can see from Fig. 10, the
model
experiences rather
weak fluctuations in the surface luminosity, reflecting the alternation
between the H- and He-burning shells in providing nuclear energy
to the star. The long-term behaviour of the surface properties indicates
the model is evolving down on its Hayashi track, i.e.
to decreasing luminosity and increasing effective temperature.
As already mentioned, the CNO surface abundance is zero, so that
the H-burning shell must itself produce the necessary 12C to
sustain the CNO-cycle. The latter is actually the dominant energy
source during most part of the AGB evolution followed
by our calculations (bottom-left panel of Fig. 10).
![]() |
Figure 10:
Evolutionary properties during the first thermal instabilities of the
He-burning shell, for two Z=0 AGB models with initial masses of
![]() ![]() |
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The
model suffers well-defined He-shell flashes,
which drive the features usually found in the light curves
of AGB stars.
Namely, over a pulse cycle, the luminosity undergoes the
initial post-flash rapid dip and peak, which is followed by the
low-luminosity dip, and the final recovering
of the luminosity level that characterizes the quiescent shell
H-burning (see, for instance, Boothroyd & Sackmann 1988).
In this case, the long-term evolution is characterised by the increase
of the surface luminosity and slight decrease of the effective
temperature, i.e. the model is climbing on the AGB.
The envelope CNO abundance (by mass) left after the completion of the second
dredge-up is
.
As for the 2.5
model, the CNO cycle provides
most of the stellar energy during the stages of quiescent H-burning
(bottom-right panel of Fig. 10).
However, it should be noticed that a fraction of the CNO nuclei involved
as catalysts in the CNO cycle are not brought up to
the surface by the second dredge-up, but are synthesized
in situ by the H-burning shell via the
process.
This can be seen by comparing the total CNO abundance (by number)
in the envelope,
,
and that, for instance, at the point of maximum nuclear efficiency
within the H-burning shell,
.
If there were not production of 12C inside the shell,
we should have
,
whereas
we find
,
with
ranging from
2 10-8 to
6.5 10-8 at
subsequent stages immediately preceding a thermal pulse.
Let us now check whether our results can be recovered
within the Fujimoto scheme.
The
and
models have core masses
of about
and
,
and
in the envelope equal to 0 and
4.198 10-9, respectively.
According to Fujimoto et al. (1984)
the lowest mass model is expected to suffer thermal pulses,
since
(regardless of
),
whereas the most massive one should not, since it is located
inside the region of the plane
where thermal pulses are prohibited (see their Fig. 5).
Thus, our results would seem to disagree with the Fujimoto scheme.
This is perhaps not surprising as the predictions by Fujimoto
et al. (1984) are derived from a semi-analytical method, in which the
input physics - adopted for the integrations of deep envelope
structures - may differ from those presently used in our stellar code,
as well as for other model prescriptions,
such as the initial helium abundance (Y=0.25 in Fujimoto et al.,
Y=0.23 in our models). Other differences may also arise from
the use of analytical approximations in the Fujimoto procedure.
However,
even if our results do not strictly obey the quantitative thresholds
in
and
as pointed out by Fujimoto et al. (1984),
a qualitative agreement in the behaviour of AGB models with initial
zero metallicity can be found.
In particular, we confirm the prediction that the occurrence of thermal
pulses in massive AGB models is conditional on the pre-existence of
some CNO abundance in the envelope, that in the case of our
model is supplied by the second dredge-up during the E-AGB.
This finding is also supported by the evolutionary calculations
of zero-metallicity intermediate-mass stars
by Chieffi & Tornambé (1984) and Domínguez et al. (1999).
The strength of thermal instabilities should also depend on
the
value in the envelope, as this quantity controls the
degree of thermal coupling between the H- and He-burning shell.
In models with initial zero metallicity, we expect that
the more
is dredged-up,
the less 12C is synthesised by the 3-
reaction in H-shell during a interpulse period,
hence the stronger the He-shell flash should be.
![]() |
Figure 11:
Masses of the He core
(
![]() ![]() |
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These quantities represent fundamental information
for both hydrodynamic calculations of the explosion event (see, for instance,
Umeda et al. 2000), and synthetic derivation of the
supernova yields (see the procedure described by Portinari et al. 1998).
![]() |
![]() |
![]() |
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7.0 | 1.09 | 0.98 | 0.715 |
8.0 | 1.66 | 1.13 | 0.688 |
8.3 | 1.75 | 1.18 | 0.677 |
9.0 | 1.96 | 1.29 | 0.658 |
9.50 | 2.11 | 1.38 | 0.643 |
10.0 | 2.25 | 1.47 | 0.630 |
12.0 | 2.90 | 1.92 | 0.574 |
15.0 | 4.01 | 2.73 | 0.506 |
20.0 | 6.13 | 4.40 | 0.422 |
30.0 | 10.79 | 8.17 | 0.330 |
50.0 | 21.25 | 17.23 | 0.234 |
70.0 | 32.10 | 27.39 | 0.187 |
100.0 | 48.73 | 42.84 | 0.149 |
Comparing the results for Z=0 with those for other metallicities,
we can notice that: i) for initial masses
25-30
the core masses are almost independent of metallicity; ii) the major
differences show up at higher masses, being largely determined by the
effect of mass loss that is expected to be less
efficient at lower Z (and not applied to the Z=0 models);
iii) in the case of the Z=0 models
evolved at constant mass, both
and
follow a nearly linear relation with the stellar mass.
In Fig. 12 we compare the evolutionary paths in the H-R diagram
for selected models of different initial masses with the
loci obtained from Eq. (5), i.e.
calculating the Eddington luminosity of the photosphere for each
value of the effective temperature along the evolutionary sequence.
![]() |
Figure 12:
Evolutionary tracks in the H-R diagram (thick lines) for selected
massive models with initial masses (in ![]() ![]() |
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![]() |
Figure 13:
Effective acceleration at the photosphere, defined as
the difference between the inward gravitational acceleration,
![]() ![]() ![]() |
Open with DEXTER |
From this simple test, it turns out that,
in contrast with the
model,
the
and
models
may achieve super-Eddington luminosities
(typically at
)
on their way to the Hayashi line,
towards the end of the He-burning phase. Correspondingly,
the effective acceleration
at the surface becomes negative, due to
the increase of the surface radiative opacity (see Fig. 13).
From this point onward the predicted evolution should not be considered reliable, as the adopted assumption of hydrostatic equilibrium does not hold any longer for the outermost layers. Likely, these stars would start losing mass from their surface.
For the sake of simplicity, we do not attempt to include any prescription for mass-loss driven by super-Eddington luminosities in the most massive stars in this work. This, and other stellar wind driving mechanisms, will be analysed in more detail in a future paper dedicated to the evolution of zero-metallicity stars with mass loss (Marigo et al. 2001, in preparation).
We will now discuss the changes in the surface chemical composition due to convective dredge-up episodes. Usually these result from the penetration of the convective envelope into stellar layers which have previously been modified by nucleosynthetic processes. In this respect it is relevant to recall that the deepest inward extension of envelope convection occurs as a star reaches close to its Hayashi line. Keeping this in mind and looking Table 3 we can make the following considerations.
Between the end of the H-burning phase and the central He-ignition,
the only models that evolve to the red part
of the H-R diagram and settle onto the giant branches
(closely approximating the location of the Hayashi tracks) are those
of low-mass (
).
At this stage envelope convection starts to proceed inward, but
it does not extend significantly into the chemical profile
left by the radiative burning core.
As can be noticed
in Table 3, this first dredge-up episode in low-mass stars
changes the surface chemical abundance of helium
by negligible amounts.
For more massive models (
)
the onset of the He-burning phase takes place far away from
the Hayashi tracks. As a consequence,
the first notable feature of zero-metallicity models
is that the first dredge-up essentially does not take place.
As far as the second dredge-up is concerned, we can notice from
Fig. 2 that
after central helium exhaustion the evolution
for all models with
is characterised by a redward excursion in the H-R diagram.
However, the minimum value of the effective temperature attained -
either at the start of the AGB phase or at central carbon ignition -
has a non-monotonic dependence on the stellar mass.
Specifically, stars with initial masses
are able to
approach their Hayashi tracks, then becoming red giants.
In this mass range the second dredge-up is found to
actually occur in models with
,
the envelope penetration being insufficient at lower masses.
Moreover, we can notice from Table 3 that for
only helium increases
its abundance in the envelope, whereas for
some primary CNO
is also dredged-up to the surface.
It is worth remarking that further changes in the surface chemical abundances may be caused by the occurrence of the third dredge-up at thermal pulses (if they take place, see Sect. 6.4) during the TP-AGB phase (see, for instance, Domínguez et al. 1999). Actually, in our calculations we do not find any evidence of this process, but this is more likely due to the small number of He-shell flashes that have been followed for each evolutionary sequence.
More massive stars (
)
met the necessary conditions for carbon ignition in the core
at higher effective temperatures, i.e.
before entering the giant branch phase. As expected, these
models do not suffer any dredge-up episode, i.e. their surface
layers preserve the original composition.
Finally, stars with
settle on their Hayashi tracks already during the core He-burning
phase, and there they remain until the onset of carbon burning.
The second dredge-up takes place enriching
the surface composition both in helium and CNO elements
(see Table 3).
As already discussed in Sect. 7.2,
the results for these massive and luminous models -
made to evolve at constant mass - should be taken with some caution.
Our analysis has in fact revealed that these models may become
dynamically unstable against
radiation pressure towards the end of core helium burning.
From that point, stellar winds driven by super-Eddington luminosities
will indirectly affect the surface chemical composition by
i) reducing the stellar mass, and ii) influencing the efficiency
of the second dredge-up. To this regard,
a more careful study is going to be carried
out in a forthcoming work (Marigo et al. 2001, in preparation).
In conclusion, owing to the
particular evolutionary features of
the zero-metallicity stars, we find
that significant changes in the chemical
composition of the envelope - during the evolutionary phases
here considered - may occur within two well defined mass ranges, i.e.
,
and
.
Our results are fully illustrated in Fig. 14.
![]() |
Figure 14: Surface abundances of 4He and CNO elements as a function of the stellar mass, for the last computed model |
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Another peculiar feature of low-mass zero-metallicity stars
(with
)
is the development of loops in
the H-R diagram, corresponding to the first activation of the CNO cycle
as soon as some primary 12C is produced by
the 3-
reaction (see Sect. 5.4).
Unfortunately, the probability of detecting such loops
in a primordial simple stellar population is negligible,
as this feature is quite short-lived (a few 106 yr),
hence scarcely populated.
Finally, distinctive observational properties of
low-mass zero-metallicity stars may be derived from
their loci in the HR-diagram: they are
slightly hotter than the Pop-II stars of even the lowest
metallicity.
For instance, Cassisi et al. (1996) find that most of
the evolution of
Z=10-10 HB stars occurs to the left
(higher
)
of the RR Lyrae instability strip, thus implying a very
low probability of finding RR Lyrae pulsators among an hypothetical
population of these stars.
However, we notice that high effective temperatures are also met
at the turn-offs of low-mass stars. This might open the possibility
that, at ages >10 Gyr, the main-sequence of Z=0 populations
falls inside the faintest part of the instability strip
which contains (in order of decreasing luminosity)
three classes of variables, namely:
Cepheids, RR Lyrae and Scuti stars. In such a case, the possibility
arises that a large number of short-period, low-amplitude pulsators
might be present among Pop-III stars.
![]() |
Figure 15:
Instability strips in the
![]() |
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To investigate this possibility, the results of a preliminary analysis
of pulsation and evolution of Z = 0 stars are shown in
Figs. 15 and 16.
We have made use of the large data set of theoretical models of Cepheid
variables calculated by Chiosi et al. (1993), corresponding to luminosities
,
various values of the stellar mass
(
), and different metal and helium
abundances (Z= 0.02, 0.008, 0.004; Y=0.23, 0.24).
This data set has been extended towards lower luminosities
by means of a preliminary stability analysis, specifically performed
for zero-metallicity models with
.
The linear nonadiabatic pulsation calculations have been carried out
with the aid of the same code as in
Chiosi et al. (1993; and references therein), but incorporating the
low temperature opacities of Alexander & Ferguson (1994).
The new calculations cover the region of the HR diagram
delimited by
and
.
Combining the earlier pulsation models with the
new ones, we are now able to determine
the edges of the instability strip
from
down to
,
as illustrated in Fig. 15.
Analogously to Chiosi et al. (1993), we have also derived several
parameterised fitting formulas (given in Table 5)
as a function of helium Y and metal content Z.
They describe, for different pulsation modes, the location of the instability
strips and other relevant relations between period, luminosity, colour,
and mass. The instability strips displayed in Fig. 15
are obtained from the aforementioned analytical relations
for the red and blue edge setting Z=0 and Y=0.23.
The main interesting points are the following.
At luminosities
the instability strip
crosses the evolutionary tracks of primordial stars which should
already be dead and hence undetectable.
Towards lower luminosities the
theoretical location of the instability strip
is consistent with the predictions
for RR Lyrae variables according to Bono et al. (1995)
for fixed helium and metal content (Y=0.24, Z=0.001).
The shaded region corresponds
to the maximum extension of the strip (to the red and to the blue),
combining the results of Bono et al. (1995)
for both fundamental and first overtone pulsators.
We can notice that the RR Lyrae instability strip is almost
twice as extended to the red as our instability strip
at the same luminosities.
It is worth remarking here that,
whereas the blue edge is well determined by our stability analysis,
the red edge is not (for more details see Chiosi et al. 1993).
As a conservative approach, we define the
red edge to coincide with the maximum of the growth rate.
In this way, we are certainly underestimating the width of the
instability strip.
Anyhow, at the typical luminosities of RR Lyrae, the instability strip for the Z=0 models would cover the red portions of the HBs described by the zero-metallicity low-mass tracks plotted in Fig. 15 (see Cassisi et al. 1996, for a similar result). In other words, pulsation should set in during a period of very rapid evolution, i.e. only after these stars have burnt most of their central helium in the bluer regions close to the ZAHB. This again would make the observability almost out of reach.
Finally, going to even fainter luminosities (
)
it turns out that zero-metallicity
models with mass as low as
are expected to exhibit
pulsational properties during the main-sequence phase.
At the luminosities and effective temperatures under consideration
(see Fig. 15),
the pulsation periods can range from about 0.6 to 0.03 days.
The reliability of these provisional results is
supported to some extent by the fact that the location
on the H-R diagram and the period determinations
are consistent with the observational data of
Scuti stars -
a rich class of faint variables with
amplitudes ranging from hardly detectable to several tenths of a
magnitude - which are usually present among the Blue Stragglers
of many Galactic globular clusters.
A sample of observed
Scuti stars
(Nemec & Mateo 1990; McNamara 2000)
is shown in the MV versus B-V diagram
of Fig. 16, and compared to the theoretical
instability strip.
It is interesting to notice that, applying the fundamental-mode
period-mass-luminosity-colour relation presented in Table 5,
to the
Scuti data (Nemec & Mateo 1990; McNamara 2000)
and assuming Z=0.001,
we can immediately estimate the pulsational
masses for these variables, with a typical range between
1.2 and
.
![]() |
Figure 16:
Instability strips in the
MV - (B-V) diagram,
with the same notation as in Fig. 15.
Isochrones of ages 15.8 Gyr and 10.0 Gyr are shown
for Z=0 (continuous line) and Z=10-4 (from
Girardi et al. 1996; dot-dashed line).
Points correspond to observed ![]() |
Open with DEXTER |
A direct comparison with observations is presented in
Fig. 16, showing isochrones representative
of Pop-II (Z=10-4) and Pop-III
(Z=0; see the Appendix) stellar populations, for two
values of the age, i.e. 10.0 Gry and 15.8 Gyr.
These latter define the best range for the age of Galactic
globular clusters, as recently reviewed by Carretta et al. (2000).
We can see that in none of the four
cases displayed, the
turn-off region is located inside the instability strip,
except for the Z=0, 10.0 Gyr old isochrone,
which marginally intersects the red edge of the strip.
However, we note that,
since our calculations probably underestimate the red extension
of the instability strip
(see earlier in this section),
there is still some possibility that the Z=0 isochrones
for ages within the range under consideration
will enter significantly into the red part of the instability strip.
If this is indeed the case, then it implies that a
large number of primordial pulsators could be
present in the present era. Such stars would presumably be
similar to Scuti variables but with lower masses
at a given colour (since Z = 0 stars near the
main-sequence are bluer and more luminous than solar metallicity
stars of the same mass). They would also have longer periods
at a given luminosity and colour (since pulsation period varies
approximately as M-0.5).
Will we ever detect them?
Mode | A | B | C | D | E | F | G | H | I |
Blue Edge:
![]() |
|||||||||
P0 | 3.928 | 2.985E-2 | -5.021E-1 | 2.986E-2 | -5.130E-2 | -1.311E-3 | 5.915E-3 | -2.319E-2 | |
P1 | 3.962 | -2.0117E-1 | -6.903E-1 | 8.774E-2 | -1.298E-2 | -1.945E-2 | 3.073E-2 | -2.658E-2 | |
P2 | 3.936 | 4.816E-2 | -8.542E-1 | 8.414E-2 | -3.018E-2 | -1.515E-2 | 1.758E-2 | 6.488E-2 | |
Red Edge: log
![]() |
|||||||||
P0 | 3.879 | 5.948E-2 | 2.274E-1 | 6.708E-2 | -4.646E-2 | -3.762E-3 | -4.628E-4 | -1.511E-1 | |
P1 | 3.901 | -7.233E-3 | 5.884E-2 | 6.731E-2 | -4.071E-2 | -6.632E-3 | 6.754E-3 | -1.177E-1 | |
P2 | 3.876 | 1.134E-1 | -5.327E-1 | 2.224E-2 | -3.824E-2 | -3.246E-3 | 2.557E-3 | 3.747E-2 | |
Blue Edge:
![]() |
|||||||||
P0 | 3.102E-1 | -1.925E-1 | 5.349 | -5.217E-2 | 4.720E-2 | -2.590E-1 | 6.541E-3 | -1.019E-2 | 1.516E-1 |
P1 | -2.215E-1 | 1.653E | 8.027 | -3.594E-1 | 1.406 | 2.767 | -2.355E-2 | 1.524E-1 | 7.409E-1 |
P2 | 1.062E-1 | 3.175E-1 | 4.482 | -2.750E-1 | 1.004 | 2.314 | -2.932E-2 | 1.664E-1 | 7.820E-1 |
Blue edge:
![]() |
|||||||||
P0 | 5.029E-1 | -4.736E-1 | 2.331 | -3.438E-2 | -9.280E-2 | -1.857E-1 | 6.263E-3 | -2.051E-2 | 3.047E-2 |
P1 | 8.104E-2 | 8.586E-1 | 3.731 | -2.736E-1 | 9.622E-1 | 1.408 | -1.644E-2 | 1.071E-1 | 3.687E-1 |
P2 | 3.184E-1 | -1.931E-1 | 1.809 | -2.387E-1 | 7.655E-1 | 2.243 | -2.424E-2 | 1.377E-1 | 6.635E-1 |
Red edge:
![]() |
|||||||||
P0 | 3.349E-1 | 1.407E-1 | 4.187 | -1.405E-1 | 3.711E-1 | -1.638 | -5.235E-3 | 5.933E-2 | 3.714E-1 |
P1 | 1.341E-1 | 8.584E-1 | 6.649 | -2.564E-1 | 8.899E-1 | 1.708 | -2.054E-2 | 1.237E-1 | 9.079E-1 |
P2 | 3.258E-1 | 7.263E-2 | 4.307 | -1.777E-1 | 5.641E-1 | 2.764E-1 | -1.799E-2 | 9.780E-2 | 2.353E-1 |
Red edge:
![]() |
|||||||||
P0 | 5.525E-1 | -1.107E-1 | 3.013E-1 | -9.441E-2 | 1.655E-1 | -1.415 | -1.833E-3 | 3.032E-2 | 1.165E-1 |
P1 | 3.827E-1 | 4.601E-1 | 2.334 | -1.970E-1 | 6.056E-1 | 1.088 | -1.605E-2 | 8.835E-2 | 5.258E-1 |
P2 | 5.315E-1 | -1.852E-1 | 6.837E-1 | -1.495E-1 | 4.114E-1 | 1.909E-2 | -1.790E-2 | 8.645E-2 | 2.669E-2 |
![]() |
|||||||||
P0 | 12.150 | 0.659 | 3.000E-2 | -0.723 | -3.567 | ||||
P1 | 10.256 | 0.683 | 2.269E-2 | -0.671 | -3.112 | ||||
P2 | 10.002 | 0.692 | 1.978E-2 | -0.685 | -3.070 | ||||
![]() |
|||||||||
P0 | -6.656E-1 | -3.771E-1 | -7.788 | -3.411E-1 | 4.206E-1 | 6.834E-1 | 5.034E+2 | 1.095E-1 | -2.927E+2 |
P1 | -6.072E-1 | -5.719E-1 | -8.367 | -3.254E-1 | 2.656E-1 | 8.998E-1 | 4.089E+2 | -7.956E-2 | -1.571E+2 |
P2 | -5.760E-1 | -7.067E-1 | -1.016E+1 | -3.099E-1 | -4.138E-1 | 1.035 | 1.092E+3 | -1.451E-1 | -1.345E+3 |
![]() |
|||||||||
P0 | -6.864E-1 | -3.636E-1 | -2.742 | -3.347E-1 | -8.353E-1 | 2.243E-1 | 1.317E+2 | 5.397E-1 | -1.859E+2 |
P1 | -6.103E-1 | -5.764E-1 | -5.814 | -3.159E-1 | -1.017 | 4.879E-1 | 2.659E+2 | 3.382E-1 | -2.217E+2 |
P2 | -5.661E-1 | -7.197E-1 | -6.262 | -3.007E-1 | -1.382 | 6.468E-1 | 6.420E+2 | 1.995E-1 | -8.657E+2 |
Acknowledgements
It is a pleasure to thank B. Salasnich for his help with the evolutionary code, and A. Weiss for useful conversations on Pop-III stars. A sincere thank goes to the referee whose comments contributed to improving the presentation of the work. This study is funded by the Italian Ministry of University, Scientific Research and Technology (MURST) under contract "Formation and evolution of Galaxies'', No. 9802192401.
The data tables for the present evolutionary tracks are available only in electronic format, either upon request to the authors, or by accessing a WWW site containing a complete data-base (including additional data and the future extensions) at http://pleiadi.pd.astro.it.
For each evolutionary track, the corresponding data file is organised into 21 columns with the following information:
age/yr
: stellar age in yr;
logL
: logarithm of surface luminosity (in solar units),
logTef
: logarithm of effective temperature (in K),
grav
: logarithm of surface gravity (in cgs units);
logTc
: logarithm of central temperature (in K);
logrho
: logarithm of central density (in cgs units);
Xc,Yc
: mass fraction of either hydrogen (up to the
central H-exhaustion) or helium (later stages) in the stellar centre;
Xc_C
: mass fraction of carbon in the stellar centre;
Xc_O
: mass fraction of oxygen in the stellar centre;
Q_conv
: fractional mass of the convective core;
Q_disc
: fractional mass of the first mesh point where
the chemical composition differs from the surface value;
L_H/L
: the luminosity
provided by H-burning reactions as a fraction of the surface luminosity;
Q1_H
: fractional mass of the inner border of the
H-rich region;
Q2_H
: fractional mass of the outer border of the
H-burning region;
L_He/L
: the luminosity
provided by He-burning reactions as a fraction of
the surface luminosity;
Q1_He
: fractional mass of the inner border of the
He-burning region;
Q2_He
: fractional mass of the outer border of the
He-burning region;
L_C/L
: the luminosity
provided by C-burning reactions as a fraction of the surface luminosity;
L_nu/L
: the luminosity
lost by neutrinos (hence negative) as a fraction of
the surface luminosity;
Q_Tmax
: fractional mass of the point with the highest
temperature inside the star;
stage
: label indicating particular evolutionary stages.
ZAMS
or
ZAHB
), local maxima and minima of L and
Te-M
,
Te-m
, L-M
, and L-m
), the exhaustion of central
hydrogen (Xc=0
) and helium (Yc=0
), the
lowest L and highest
Bhe
and LpT
, respectively),
the base and tip of the first ascent of the red giant branch (Brg
and Tip
, respectively), the maximum L immediately preceding a
thermal pulse (1tp
), and the onset of C-burning (Cb
).
These stages delimit characteristic evolutionary phases, and can be useful
for the derivation of physical quantities (as e.g. typical lifetimes) as a
function of either mass or metallicity.
Notice that some of these evolutionary stages may be absent from
particular tracks, depending on the precise value of stellar mass
and metallicity.
From the tracks presented in this paper, we have constructed isochrones
adopting the same algorithm of "equivalent evolutionary points''
as used in Bertelli et al. (1994) and Girardi et al. (2000).
The initial point of each isochrone is a 0.7
model in the
main sequence. The terminal stage of the isochrones corresponds either to
stars at the beginning of the TP-AGB for ages
108 yr,
or to stars at C-ignition for
108 yr.
The youngest isochrones (up to 6 Myr) are basically limited to the
main sequence stages, since at these ages the inclusion of more
evolved stars would require tracks for
,
that will
be presented in a forthcoming paper.
Complete tables with the isochrones can be obtained
upon request to the authors, or through the WWW site
http://pleiadi.pd.astro.it.
They cover the complete age range from about 2 Myr to 16 Gyr
(
). Isochrones are provided
at
intervals; this means that any two consecutive
isochrones differ by only 12 percent in their ages.
Figure B.1 shows some of these isochrones
on the HRD, sampled at age intervals of
.
Theoretical luminosities and effective temperatures along the isochrones
are translated to magnitudes and colors by using extensive tabulations of
bolometric corrections and colors, as in Bertelli et al. (1994).
The tabulations were obtained from convolving the spectral
energy distributions contained in the
library of stellar spectra of
Kurucz (1992) with the response function of several broad-band filters.
The response functions are from Buser & Kurucz
(1978) for the UBV pass-bands, from Bessell (1990) for the R and I
Cousins, and finally from Bessell & Brett (1988) for the JHK ones.
For each isochrone table, the layout is as follows:
A header presents the basic information about the age and metallicity
of each isochrone. Column 1 presents the logarithm of the age in yr;
Cols. 2 and 3 the initial and actual stellar masses,
respectively. We recall that the initial mass is the useful quantity
for population synthesis calculations, since together with the initial
mass function it determines the relative number of stars in different
sections of the isochrones. Then follow the logarithms of surface
luminosity (Col. 4), effective temperature (Col. 5), and surface
gravity (Col. 6). From Cols. 7 to 15, we have the sequence of
absolute magnitudes, starting with the bolometric one and following
those in the UBVRIJHK pass-bands. In the last Col. (16), the
indefinite integral over the initial mass M of the initial mass
function (IMF) by number, i.e.
![]() |
(B.1) |
We also provide summary tables containing basic information for the most significant stages along the isochrones. The evolutionary stages listed are, in sequence:
TO
: the turn-off point, i.e.
the point of highest
Te-m
and Te-M
signal
the coldest and hottest points, respectively, of stars in the
overall contraction phase at the end of core-H burning;
RGBb
and RGBt
: the base and tip of the RGB,
respectively;
BHeb
: the beginning of the CHeB phase;
Te-m
and Te-M
signal
the coldest and hottest points, respectively, for CHeB stars;
For the youngest isochrones, Te-M
represents the maximum
extension of the Cepheid loop;
EHeb
: the end of the CHeB phase;
L-M
and L-m
limit the luminosity range of early-AGB stars; this interval
corresponds to the clump of early-AGB stars in colour-magnitude
diagrams;
1TP
: the beginning of the thermally pulsing
AGB phase;
Cb
: the stage of C-ignition in the cases the
AGB phase do not occur.