A&A 374, 629-637 (2001)
DOI: 10.1051/0004-6361:20010725
P. Lundqvist1 - C. Kozma1 - J. Sollerman1,2 - C. Fransson1
1 - Stockholm Observatory, 133 36 Saltsjöbaden, Sweden
2 - European Southern Observatory, Karl-Schwarzschild-Strasse 2,
85748 Garching bei München, Germany
Received 2 March 2001 / Accepted 17 May 2001
Abstract
ISO/SWS observations of SN 1987A on day 3425 show no emission in
[Fe I] 24.05
and [Fe II] 25.99
down to the limits
0.39 Jy and
0.64 Jy, respectively. Assuming a homogeneous
distribution of 44Ti inside
and negligible dust cooling,
we have made time dependent theoretical
models to estimate an upper limit on the mass of ejected 44Ti.
Assessing various uncertainties of the model, and checking the late optical
emission it predicts, we obtain an upper limit of
.
This is lower than in our previous estimate using other ISO data,
and we compare our new
result with other models for the late emission, as well as with expected yields
from explosion models. We also show that steady-state models for the optical
emission are likely to overestimate the mass of ejected 44Ti.
The low limit we find for the mass of ejected 44Ti could be higher if dust
cooling is important. A direct check on this is provided by the
gamma-ray emission at 1.157 Mev as a result of the radioactive decay
of 44Ti.
Key words: supernova: individual: SN 1987A - nucleosynthesis - supernovae: general
While
and
are rather accurately known for SN 1987A,
with
(e.g., Suntzeff & Bouchet 1990),
and
(Fransson & Kozma 1993, and references
therein), the situation is more uncertain for 44Ti. To some extent this
has been due to a poorly known decay time of this isotope, although this has
recently improved; 44Ti decays first to 44Sc, on a time scale
of
years (Ahmad et al. 1998; Görres et al. 1998; see also
Norman et al. 1998 who obtain
years), and then quickly further
to 44Ca.
Based on time dependent modeling, K00 estimates
that
best explains the evolution
of the optical broad-band photometry of SN 1987A for
days.
This is consistent with Chugai et al. (1997), who
find
-
from the optical line emission
at 2875 days, and Mochizuki & Kumagai (1998)
who obtain the same result from a light curve analysis (see also Nagataki
2000). However, as is shown by K00, both broad-band photometry and the
emission modeled by Chugai et al. (1997) constitute only a minor fraction
of the total emission put out by the supernova. At these epochs most of the
emission from the supernova instead comes out in the far infrared (IR)
in a few iron lines, most notably [Fe II] 25.99
.
This makes bolometric corrections to the late
optical data rather uncertain, and a more direct way to
measure the content of
is to measure the flux in the far-IR
lines.
In a recent study Lundqvist et al. (1999; henceforth L99) analyze
Infrared Space Observatory (ISO; Kessler et al. 1996) data obtained mainly
at t = 3999 days. L99 couple these observations to time dependent
model calculations similar to those in K00,
and from the absence of iron line emission in the
ISO spectra they derive
.
This limit is
consistent with the results of Chugai et al. (1997), Mochizuki & Kumagai
(1998) and K00. A different limit was obtained by Borkowski et al. (1997)
who in their preliminary analysis of ISO data from t = 3425 days
derive
.
Such a low titanium mass would have
important consequences for models of the explosion. Here we analyze the data
of Borkowski et al. in the same way as was done for the 3999 day data in
L99, however, expanding the discussion on the uncertainties of our
method. In particular, we discuss the uncertainty of the temperature,
which is important since the far-IR lines are collisionally excited, and
their excitation energy (e.g.,
K for the 26
line)
is much higher than the gas temperature. Because the emission at this
epoch is powered mainly by positrons from the decay of 44Ti, we also
discuss the uncertainty of the line fluxes due to the efficiency of the
trapping of positrons.
We have used the ISO data discussed in Borkowski et al. (1997). These data
were obtained on 10 July, 1996 (i.e., on day 3425 of SN 1987A)
in the SWS06 mode of the Short Wavelength Spectrograph (SWS; de Graauw et al.
1996). The data we
analyze span the regions
and
.
Both these ranges are covered by band 3D of SWS, and include the lines
[Fe I] 24.05
and [Fe II] 25.99
,
i.e., the two strongest lines
expected from the supernova (cf. L99). These data are superior to
those discussed in L99 as the SWS06 scans are deeper than the
SWS01 scans in L99, and include a large enough wavelength interval for
accurate continuum determination compared to the SWS02 data in L99.
The data were retrieved from the ISO archive, and have been reduced
with the off-line processing (OLP/pipeline) version 8.7. We used ISO
Spectral Analysis Package (ISAP) to flatfield and average over the scan
directions. The averaging over the detectors was done
with a resolution of R = 300, using the standard -clipping.
![]() |
Figure 1:
ISO SWS/AOT6 spectra of SN 1987A on day 3425. The spectra
were reduced using the pipeline software (see text). The bin size is
![]() ![]() ![]() ![]() ![]() ![]() ![]() ![]() |
Open with DEXTER |
In Fig. 1 we present fully reduced spectra of SN 1987A for the two
wavelength regions. Although the nominal instrumental resolution of the SWS
spectra is
we have averaged the spectra with a bin size
of R = 300, corresponding to
.
This is sufficient
to resolve the lines since they could extend to well over
.
However, we detected neither [Fe I] 24.05
nor [Fe II] 25.99
in
the spectra. We estimated the rms in the data with a zero-order
baseline fit in the regions 23.80-24.80
and 25.60-26.40
,
and found 0.11 Jy and 0.19 Jy, respectively.
The exact rms-values vary somewhat with the bin size used in the
averaging routine. Reducing the bin size to R = 500 gives
the limits 0.13 Jy and 0.21 Jy, respectively. In order to
derive a conservative limit on the 44Ti mass we have adopted these
values. A 3
limit for the
line then becomes
0.64 Jy, while for the
line it is lower,
0.39 Jy. The bin
size is small enough so that we can use the limits for the peak of the
expected line profiles in our estimate of the 44Ti mass.
The thermal and ionization balances are solved time-dependently,
as are also the level populations of the most important ions.
A total of 6400 lines are included in the calculations.
The radioactive isoptopes included are 56Ni, 57Ni, and 44Ti,
and we calculate the energy deposition of gamma-rays and positrons solving the
Spencer-Fano equation (see Kozma & Fransson 1992). We assume that the
positrons deposit their energy locally.
We include
of 56Ni and
of 57Ni (Sect. 1). In L99 three values of
were
tested:
and
.
Here we extend
this to include three new values:
and
.
Although we follow the complete evolution of the
supernova after 150 days, we concentrate our discussion mainly on 3425
and 3999 days, i.e., the epochs of the ISO observations.
In L99 an e-folding time of 78 years was used for 44Ti, as this was
thought to represent a mean value from experiments. More accurate measurements
of the decay time (Ahmad et al. 1998; Görres et al. 1998; Norman et al. 1998)
became available while the analysis of L99 was completed, and the results of
L99 were corrected accordingly before publishing. Here, we use the
more accurate e-folding time (
years) from the outset.
The explosion model we use is the same as in L99
and Kozma & Fransson (1998a). That is, we take the abundances from the 10H
model (Woosley & Weaver 1986; Woosley 1988), but distribute the
elements so that hydrogen is mixed into the core.
Spherically symmetric geometry is assumed, and the iron-rich core extends out
to
,
outside of which (out to
)
a hydrogen envelope
is attached.
We use the Sobolev approximation for the line transfer. This is a good approximation for well-separated lines in an expanding medium, but is poorer when lines overlap. This is actually the case in the UV, especially at earlier epochs. The overlap leads to UV-scattering which affects the UV-field within the ejecta. As in L99, we have studied two extreme cases to test this effect: photoionization as in the original model, and simply switching off the photoionization caused by the UV-field. From the results of L99 we, however, do not expect the UV-field to be the dominant source of uncertainty in our models. Another drawback of the Sobolev approximation is that we do not account for line fluorescence in which UV lines are split into optical and IR lines. To treat this accurately is beyond the scope of this paper, but is treated in detail in a forthcoming paper.
Since L99 we have updated the collision strengths for the [Fe II]
transition. The collision strength,
,
is given by Zhang & Pradhan
(1995) as a function of temperature. The value of
at
K
is set to
(A. Pradhan, private communication).
We note, however, that there can still be uncertainty in this result
as a fully relativistic calculation is needed to accurately calculate
the collision strength at these low temperatures (M. Bautista, private
communication).
|
Photoion.b |
![]() |
![]() |
|
yes | 0.047 | 0.095 |
![]() |
no | 0.051 | 0.074 |
![]() |
yes | 0.072 | 0.17 |
![]() |
no | 0.079 | 0.14 |
![]() |
yes | 0.11 | 0.34 |
![]() |
no | 0.12 | 0.29 |
![]() |
yes | 0.17 | 0.61 |
![]() |
no | 0.18 | 0.53 |
![]() |
yes | 0.26 | 1.37 |
![]() |
no | 0.31 | 1.18 |
![]() |
yes | 0.40 | 2.96 |
![]() |
no | 0.54 | 2.53 |
a Distance = 50 kpc. Homogeneous distribution of emitting gas
inside
![]() The resulting shape of the line profile is seen Fig. 1. The flux values are the peak values of the lines assuming this profile. b Indicates whether or not photoionization has been included in the calculation. (See Sect. 3.1). |
The line fluxes in Table 1 are for a distance to the supernova of 50 kpc, and
assuming a homogeneous distribution of emitting gas inside
.
The FWHM velocity for the resulting profile (shown in Fig. 1)
is
,
which is close to what Haas et al. (1990)
observed for [Fe II] 17.94
at
400 days after the
explosion,
.
![]() |
Figure 1:
Modeled fluxes of [Fe I] 24.05
![]() ![]() ![]() |
Open with DEXTER |
![]() |
Figure 2:
Fraction of iron in Fe II,
![]() ![]() |
Open with DEXTER |
In L99 it was discussed how the fluxes in [Fe I] 24.05
and
[Fe II] 25.99
(
and
,
respectively) scale
with increasing values for
at a given epoch. It was found
that
scales nearly linearly with
,
while
has a weaker dependence. The stronger dependence for
is
due to an ionization effect because
,
the relative fraction of
iron in Fe II in the iron-rich gas, increases with increasing
.
We find that this trend continues also for
.
This is shown in Figs. 2 and 3, where models
joined by dashed (solid) lines are without (with) photoionization.
Whereas
and the temperature in the iron-rich gas are only
shown for 3425 days (Fig. 3),
results for the line fluxes (Fig. 2) are shown for both 3425 and 3999
days. Note how the fluxes remain nearly constant between these two epochs,
especially in the case of
.
The only deviation from a linear log(
)
versus log(f) behavior in
Fig. 2 is for
at the highest values of
in the
models with photoionization. This is because here
(the fraction of iron in Fe I in the iron-rich gas) starts to fall
significantly below unity (cf. Fig. 3):
for
,
at 3425 days. (The same number for the case without
photoionization is
.)
The temperature increases monotonically with
(Fig. 3) and
affects
and
in the same way. This is because
the two lines have nearly the same excitation energy, and their effective
collision strengths vary only weakly (and in our models are assumed to be
constant over the small temperature regime [
115-180 K] found in the
iron-rich gas in our models).
Table 1 and Fig. 2 show that the 26
line is
stronger than the 24
line even
for the models with the lowest values of
(and thus the highest
values of
). This is because the collision strength for
[Fe II] 25.99
line is much larger than for [Fe I] 24.05
.
Table 1 and Fig. 2 also show that switching off photoionization
does not have a dramatic effect on the line fluxes. Again, the largest
difference for the fluxes is for the 24
line and at the
highest
considered. For values
,
switching
off photoionization only affects
(
)
by
11% (
22%) at 3425 days.
The small difference is simply due to the minor shift in the degree of
ionization, from Fe II to Fe I, when photoionization is switched off.
With the results in Fig. 2 it is straightforward to estimate the upper limit
on
.
It is clear that the best estimate within the framework of our
modeling comes from the 26
line on day 3425.
Using
Jy from Sect. 2.1, we find the upper
limit
with (without)
photoionization included. A conservative limit (i.e., the case when
photoionization is unimportant) is
therefore
.
In Fig. 1 we have included the expected line emission for a model
with
,
i.e., close to this limiting mass.
Although our limit is nearly a factor of four higher than that found by
Borkowski et al. (1997), using the same data, it is still substantially lower
than found by L99 from day 3999. Using our updated code, the day 3999
data only set a limit
.
This is a
factor
1.7 better than can be obtained from the 24
line on
day 3425. We will evaluate our limit from the 26
line on day 3425
in Sect. 4.1.
A source of uncertainty, not investigated in L99, is the rather unknown
rates of charge transfer for the elements in the core. In the models
in Sect. 3, no charge transfer was included between Fe and excited states
of He, as well as between atoms and ions of Fe. We have tested the importance
of including this charge transfer, using the rates suggested by Liu et al.
(1998) in a model with
and no photoionization.
Compared to the model with the same
in Sect. 3, the model with the
full charge transfer has a higher value of
by
20%,
mainly due to charge transfer between Fe I and Fe III. However, the
temperature in the model with full charge transfer is somewhat lower,
so the difference in
between the two models is
only
7%. With full charge transfer included, the estimated
would therefore be lower than in Sect. 3, but due to the uncertainty
of the charge transfer rates, and in order to derive a conservative limit
on
,
we assign a generous error of 10% in
due to
this effect.
L99 studied the uncertainty of the collision strength of the 26
line. Although we are now using a better fit to the results
of Zhang & Pradhan (1995), the collision strength at very low temperatures
is still uncertain (see Sect. 3.1). In L99 we assumed an uncertainty
of
in
due to this, and we retain this number also here.
The collisional excitation of the 24
and 26
lines is more
sensitive to the temperature (through the
term) than to the collision strength. Too high a temperature in our
models could therefore overestimate the flux in these lines, and we
would consequently underestimate
.
We can check this effect by
studying Table 1 where we see that models with
and
differ in
by a factor
of
2.2 at 3425 days. Because the temperature in
the
model is
159 K (see Fig. 3),
a lowering of the temperature by only
20% would
decrease
to the same level as in
the
model, for constant
.
However, the temperatures in our models are not free parameters but are fixed
by the heating and cooling. The energy not coming out in the 24
and 26
lines must instead come out in other lines (like the optical/IR
recombination lines) which are less sensitive to temperature, or as emission
from dust. We will make a consistency check of the optical and far-IR emission
in our models in Sect. 4.2, but we note already here that only a small
transfer of energy loss from the far-IR lines to the optical/IR recombination
lines would increase the optical and IR flux considerably, and then the models
of K00 for the optical would severely overestimate
.
We therefore believe we are not making a serious error in the derivation of
the temperature
in our models.
Dust emits a continuum which is difficult to detect, but we may for simplicity
assume that the dust emission does not interact with the gas. In that case,
dust cooling has exactly the same effect for the 24
and 26
line fluxes as just lowering
.
(See the models in Sect. 3.)
It may therefore well be that dust cooling could tap
the supernova of its [Fe II] 25.99
emission, but it is likely that it
would do so in a way which would also make the optical and near-IR lines
too weak. We do not assign an explicit error in the
estimated
due to dust cooling. (See also Sect. 4.2.)
L99 discussed uncertainties in the modeled line fluxes due to the explosion
model used. As in L99 we have used the 10H explosion model (Woosley & Weaver
1986; Woosley 1988), mixed by Kozma & Fransson (1998ab) to
give good agreement between their modeling and late time observations.
Kozma & Fransson (1998b) compared the results from this model with
a similarly mixed version of the 11E1 model (Shigeyama et al. 1988),
and found that the iron lines are insensitive to the explosion model used.
This is because the iron core mass, which is fixed by the amount of
ejected 56Ni, is the same in both models. We follow L99 and assign
a
error in
due to the choice of explosion model.
In our calculations we have assumed a local deposition of positron energy
from the radioactive decay of 44Ti. This assumption was also used
in L99. The argument is that the efficiency of trapping cannot change
substantially until day 3600 (K00), and that the most straightforward
interpretation of a near-constant efficiency of trapping over an extended
period of time is to assume full trapping. L99 assumed a
error
in
due to the uncertainty of trapping. Here we do not infer an
explicit error to the models in Sect. 3, but in Sect. 4.2 we investigate
this in greater detail in terms of a consistency check of the modeled IR and
optical fluxes. The positron deposition is discussed in Sect. 4.2
in conjunction with clumping. Even without positron leakage,
clumping of the iron-rich gas could cause an error in the estimate of
.
However, according to L99, this error is likely to be small,
5%, and
can be ignored when compared to other errors discussed above.
The combined error of
due to model approximations,
except for different clumping and positron leakage scenarios (which will
be discussed in Sect. 4.2)
is therefore
36%. With the line profile used in Fig. 2 we thus
arrive at an upper limit on
which is
.
Although the limit on
we found above in Sect. 4.1 is a factor
of
2 lower than the limit found by L99 for the 3999 data,
it is still compatible with K00 (see Sect. 1). However, the
agreement is not perfect although the same computer code has been used.
Here we make a consistency check of the modeled optical and IR emission
produced by our model. For the late optical photometry (day 3268)
we adopt the same data as used by K00 (Soderberg et al. 1999).
The observational error of these is approximately
magnitudes, as judged from the systematic error in the zero point of HST/WFPC2
photometry (Mould et al. 2000; P. Challis 2000, private communcation).
We have concentrated our consistency check on models
with
.
This is close to the upper limit found
in Sect. 4.1 from the far-IR lines when the errors have been considered. This
value of
is also in the lower range of the preferred values found
by K00 in her modeling of the optical and near-IR emission. This can be
seen from Table 2 where the modeled V, R and I magnitudes in
model M1 are all slightly fainter, by roughly half a magnitude, than
the observed, while the predicted [Fe II] 26
line emission is somewhat
too strong to be accommodated by our error analysis in Sect. 4.1.
(Model M1 is identical to the model with
and no
photoionization in Table 1.)
It is obvious that our models contain simplifications that can make us
systematically overproduce far-IR emission at the expense of optical
emission. We still have to explore fluorescence (see Sect. 3.1), and other
simplifications could be clumping and different scenarios
for the deposition of positron energy. These two effects are probably also
coupled. In Sect. 4.1 we mentioned that L99 found that the
estimated error in
due to clumping alone should be small. This is
true if one looks at the total emission coming out from the core, and if
the energy deposition is local. However, if we assume that the positron
optical depth scales linearly with density (ignoring effects due to
the magnetic field), positron leakage from the low-density component is
more likely than from the high-density one.
To test this we have, similar to what we did in L99, used a two-component model
with equal mass in the dense and the less dense components. For the density
ratio we use the value 5. From a purely geometrical point of view, positrons
produced in a high-density clump see an optical depth for absorption in the
positron emitting clump which is
3 times higher than for self-absorption
of positrons emitted in a low-density region. It is therefore more likely that
positrons from the low-density gas leak from their production site, and the
chance that they are captured in either a low-density, or high-density
component on a larger scale is about equal. From a physical point of view,
absorption of the leaking positrons is slightly biased toward absorption in
the low-density gas since this is more ionized (due to less efficient
recombination), which results in slightly higher deposition there. We have
tested three scenarios, all without photoionization and
with
:
the first (M2 and M3) is the same as in L99
with local deposition of the positron energy. The second (M4 and M5) is with
all positron deposition (i.e., from both the low and high-density gas)
occurring only in the high-density clumps. This situation is an extreme
case (cf. above), but could be possible if magnetic effects make deposition
in the high-density gas more likely than in the low-density gas. The third
(M6 and M7) scenario is for a situation where the positrons in the low-density
gas are instead assumed to leak into and deposit their energy in the the
silicon-rich gas which is macroscopically mixed with the iron-rich gas.
We have run these models both as steady state and time dependently to separate
this effect as well. We have summarized the results in
Table 2, where we show
,
,
V, R and I
at 3425 days. (The observed optical fluxes in Table 2 are for
3268 days, and not for 3425 days. The change in modeled optical flux
between these two epochs is, however, negligible.)
Model |
Time dep.b |
![]() |
![]() |
V | R | I |
M1d |
yes | 0.32 | 1.18 | 20.4 | 19.7 | 20.0 |
M2e | yes | 0.39 | 1.30 | 20.3 | 19.7 | 20.0 |
M3e | no | 0.35 | 1.27 | 20.3 | 20.4 | 20.2 |
M4f | yes | 0.47 | 1.13 | 20.1 | 19.6 | 19.9 |
M5f | no | 0.46 | 1.15 | 20.1 | 20.4 | 20.2 |
M6g | yes | 0.32 | 0.80 | 19.6 | 19.5 | 19.7 |
M7g | no | 0.30 | 0.82 | 19.6 | 20.0 | 19.9 |
From the results in Table 2 it is clear that the far-IR emission is
rather robust to density inhomogeneities in the core, as well as to the
assumption of steady state as long as the positrons are all trapped
in the iron-rich gas. Allowing for positron diffusion into the Si-rich
gas decreases the [Fe II] 25.99
emission; instead the emission in Vincreases notably, while R and I are less affected. In our most
extreme model M6 (cf. Table 2) the positron escape leads to a reduction
of
by
32%. This model actually produces 26
emission which is below the detection threshold, so this shows that for
extreme situations of clumping and positron depositions our model could
accommodate up to
in SN 1987A, and still not
produce detectable far-IR emission. However, such a scenario would
imply
at 3425 days, a factor
1.6 higher
flux in V than observed.
As the optical flux is only a minor contribution to the total emission
coming out from the supernova at this epoch, and we yet have to investigate
fluorescence to model the optical emission in detail, we cannot exclude
a scenario like this. We therefore base our upper limit on the 26
line which we believe we have modeled more accurately. To estimate a
conservative upper limit on
,
taking into account various clumping
and deposition effects, we thus arrive at an upper limit of
.
Models for the yield of 44Ti give quite different results. In particular,
in the model of Timmes et al. (1996; see also Woosley & Weaver 1995) with
a zero-age mass of
(i.e., corresponding to SN 1987A) the
mass of the initially ejected 44Ti is
,
but only
escape after fallback. This is in accord with
our upper limit in Sect. 4.1, but it should be emphasized that the variation
of
with
in Timmes et al. (1996) is complex, and for models
with
and
,
the calculated
is closer to our limit. Furthermore, the recent models of Hoffman et al.
(1999) for 15 and 25 solar mass models, indicate that the yield of 44Ti
in Woosley & Weaver (1995) may have to be increased.
The models of Thielemann et al. (1996) have larger entropy and thus more
alpha-rich freeze-out than those of Woosley & Weaver (1995). Accordingly, the
ratio
(where
is the mass of ejected 56Ni that
does not fall back) is higher. For example, in their
model
and
.
This value for
is higher than our upper limit in
Sect. 4.2. With the refinements made by Hoffman et al. (1999) to the models
of Woosley & Weaver (1995), it seems that our upper limit is lower than, or
at least close to,
in the models of both groups. The calculated
is, however, very model dependent, and sensitive to, e.g., the
explosion energy, which in the case of SN 1987A is still only known to an
accuracy of
30% (Blinnikov et al. 2000).
The models of Nagataki et al. (1997, 1998) are related to those of Thielemann
et al. (1996), although they also allow for 2-D. With no asymmetry,
Nagataki et al. obtain
for a SN 1987A-like explosion, which is a factor of
3 lower than
Thielemann et al. (1996), and could indicate the range of uncertainty of the
modeling. In the models of Nagataki et al. the yield of 44Ti quickly
increases with the degree of asymmetry. For an asymmetry of 2 between the
equator and the poles, the stronger alpha-rich freeze-out in the polar
direction increases
to
-
(for
), the range of titanium yield depending on the
form of the mass cut (Nagataki 2000). Nagataki (2000) argues that this degree
of asymmetry gives a distribution of ejected 56Ni which agrees with
observed line profiles of [Fe II] 1.26
(Spyromilio et al. 1990) and
[Fe II] 18
(Haas et al. 1990) at
400 days.
Although the models of Nagataki (2000) thus hint that asymmetry can explain the
distribution of the ejected 56Ni in SN 1987A, the need for asymmetry
to reach 44Ti in excess of
cannot be considered a strong
case, given the uncertainty in the modeling. We therefore do not
regard our upper limit on
for SN 1987A
as contradictory when compared to the absolute yield of 44Ti in,
e.g., Nagataki (2000), especially since the models of Nagataki are trimmed
to produce
in agreement with the light curve results of Mochizuki &
Kumagai (1998).
Table 2 also shows that although steady-state models can be used for the
far-IR emission from the core, they fail to reproduce the optical
broad-band emission. In particular, they underproduce R and I
magnitudes. For the cases in Table 2, the maximum errors are 0.8 and 0.3
magnitudes, respectively. This is due to the freeze-out effect
described by Fransson & Kozma (1993). Consequently, any estimate
of
based on steady-state calculations for the optical broad bands
is likely to overestimate
,
even though the structure and
atomic data are correct. It could be that the models of Mochizuki &
Kumagai (1998; see also Nagataki 2000) suffer from this, which may
explain why their lower limit,
,
almost coincides
with our upper limit, and why their lower limit is higher than the one
of K00.
Finally, we note that our upper limit on
could be higher if dust
contributes significantly to the cooling at late times. This
is not unique to our models, but affects all models trying to reproduce
the optical and IR emission. The range of
bracketed by K00
and our analysis,
would then be shifted
to a range with higher masses.
Heavy dust formation in the core
would, however, block out the line emission from the core and affect line
profiles. To test the importance of dust cooling, direct measurements
of the gamma-ray emission at 1.157 MeV, the result of the radioactive
decay of 44Ti, are needed. Future, and even more sensitive gamma-ray
instruments than INTEGRAL, are needed for such a study.
Acknowledgements
We thank Peter Challis, Nikolai Chugai and Stan Woosley for discussions. We are grateful to support from the Swedish Natural Science Research Council and the Swedish National Space Board.