A&A 426, 219-227 (2004)
DOI: 10.1051/0004-6361:20040451
R. Mauersberger1 - U. Ott2 - C. Henkel3 - J. Cernicharo4 - R. Gallino5,6
1 - Instituto de Radioastronomía Milimétrica, Avda.
Divina Pastora 7, Local 20, 18012 Granada, Spain
2 -
Max-Planck-Institut für Chemie, Becherweg 27, 55128 Mainz,
Germany
3 -
Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69,
53121 Bonn, Germany
4 -
Instituto de la Estructura de la Materia, Dept. de
Astronomía Molecular e Infrarroja, Serrano 113, 28006
Madrid, Spain
5 -
Dipartimento di Fisica Generale
dell'Università di Torino, Via Pietro Giuria 1, 10125 Torino,
Italy
6 - Centre for Stellar and Planetary Sciences, Monash
University, Melbourne 3800, Australia
Received 16 March 2004 / Accepted 21 June 2004
Abstract
The J=2-1 and 3-2 rotational lines of the rare
isotopomer C36S and the J=5-4 and 6-5 transitions of Si36S were detected in the carbon
star IRC+10216 (CW Leo). These are the first detections of 36S-bearing molecules in a star and the first spectroscopic
detection of Si36S. From a comparison of 34S- and 36S-bearing
isotopomers, the
34S/36S isotopic ratio is 107(). This value is comparable to values
in the interstellar medium of the inner Galactic disk (115)
but is smaller than the solar value of 288 (Ding et al. 2001). The increase of the
36S abundance relative to 34S only qualitatively follows
model predictions of a low mass AGB star. Quantitative agreement of the observed 34S/36S ratio
with the stellar models can be reached if the age of IRC+10216 and Galactic chemical evolution are taken into
account. Other less likely possibilities are the presence of
considerable inhomogeneities
in the interstellar medium and either
IRC+10216 or the Sun started with a peculiar 36S abundance. Other production mechanisms potentially
capable of enhancing the Galactic interstellar medium are
discussed. From the observed line density toward IRC+10216 and toward Galactic star forming regions,
we estimate the confusion limit toward those sources.
Key words: nuclear reactions, nucleosynthesis, abundances - stars: abundances - stars: AGB and post-AGB - stars: individual: IRC+10216 - ISM: abundances - radio lines: stars
Being one of the ten most abundant elements and possessing four stable isotopes, sulfur is of particular interest for such studies. Relative abundances of stable sulfur isotopes have been determined for meteorites (e.g. Gao & Thiemens 1991), the Moon (Thode & Rees 1971), the Galactic ISM (e.g. Chin et al. 1996), Cosmic Rays (Thayer 1997) and late type stars (e.g. Kahane et al. 1988). Even in external galaxies 32S/34S ratios could be measured (Mauersberger & Henkel 1989; Johansson et al. 1994). Small isotopic variations of sulfur in terrestrial, meteoric and planetary samples have been proposed to be powerful indicators of chemical, geophysical and biological processes (Canfield 2001; Farquhar & Wing 2003)
Mauersberger et al. (1996) presented the first
interstellar detections of a 36S-bearing molecule toward a
number of Galactic hot cores. While the solar 34S/36S ratio is 288, an interstellar ratio of 115()
was found. This
supported the notion that unlike other S isotopes, 36S is a
secondary-like nucleus, predominantly synthesized via the s-process
in massive stars. Most interstellar 34S/36S ratios were
determined, however, for sources with galactocentric distances
kpc. A positive 34S/36S gradient with
increasing
,
i.e. a ratio between 115 and
180 in
the solar neighborhood, is not completely ruled out by the data of
Mauersberger et al. (1996).
In order to further constrain the production site of 36S in the Galaxy, we conducted a search for rotational lines of C36S and Si36S toward the prototypical carbon star IRC+10216, where all other stable S isotopes have been detected previously (Kahane et al. 1988).
The frequencies of Si36S have not been measured in the
laboratory. However, it is straightforward to derive the
rotational constants of Si36S from those of SiS (see, e.g.,
Frum et al. 1990) using the isotopic relation for diatomic
molecules (e.g. Townes & Schawlow 1975). For Si36S, we derived
B0=8607.47 MHz and
MHz. The computed frequencies are
MHz and
MHz. The
expected relative error for the rotational constants is <10-5and the resulting expected uncertainty for the J=5-4 and 6-5 lines of Si36S is <0.5 MHz. We have checked the precision
of those relations for the different isotopes of SiS that have
been measured in the laboratory and we find that the quoted
uncertainties are rather conservative. Recently, Sanz et al.
(2003) measured the rotational constants of several
isotopes of SiS and have provided a fit, taking into account the
breakdown of the Born-Oppenheimer approximation, to the data of
all isotopes of SiS. From their constants we infer those of Si36S to be
B0=8607.495 MHz,
MHz,
MHz and
MHz, i.e, 0.3 MHz of difference with
respect to our early calculations. Hence, we are confident that
the calculated frequencies have an error well below 1 MHz.
![]() |
Figure 1:
The whole spectral range measured near the C36S J=2-1 ( lower panel) and 3-2 ( upper panel) transitions toward
IRC+10216. The channel spacing corresponds to
![]() ![]() ![]() |
As backends, we employed filterbanks with 256 channels and a
channel spacing of 1 MHz corresponding to
for the 3 mm transitions and 2.1 km s-1 for the
2 mm transitions. Two receivers were employed simultaneously,
either observing orthogonal polarizations at the same wavelength
or measuring one line at 2 mm and the other at 3 mm using just
one polarization for each wavelength.
All spectral lines were measured using the wobbling secondary mirror with a beam throw of 200'' in azimuth. The phase time was 2 s and the on-source integration of each subscan was 30 s. Subscans obtained by wobbling to the left and to the right of our source were added to eliminate baseline ripples caused by the asymmetry in the beam path (symmetric switching).
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Figure 2:
C34S and C36S J=2-1 and 3-2 spectra toward IRC+10216. The channel spacing corresponds to
![]() ![]() ![]() ![]() |
Figure 1 displays the entire spectral range measured
at the frequencies of the C36S lines. Figures 2 and 3 show the C34S, C36S, Si34S and Si36S spectra that are discussed in Sect. 5.2 in more detail. An inspection of
Fig. 1 shows that a main source of uncertainty in
the identification of weak lines and also in the definition of
spectral baselines comes from the presence of many weak line
features, in particular at 2 mm. In Appendix A, we list
the line parameters of all features together with possible
identifications. The Si36S data, showing a statistical
behavior that is similar to the C36S data, will be analyzed
elsewhere (Cernicharo et al., in prep.).
From the spectra, we subtracted baselines of first order. For the 3-2 line of C36S, almost the entire frequency range is covered with line features near or above the detection threshold making the identification of a region free of emission a difficult task. We chose spectral regions with particularly low emission for the definition of a baseline.
As it can be seen in Fig. 1
showing the total spectral range of
MHz toward
IRC+10216, we have detected a large number of spectral line
features, many of them unidentified. There are more lines at 2 mm
than at 3 mm although the rms of the data is similar. The spectral
line density in our 2 mm spectrum is 15 lines per 256 MHz. Since
each line has a width of about 14.5 MHz, about 85% of the spectral
range observed would be covered with detectable line emission if
there were no overlap. The weakest feature we identified in
Table A.1 at
mm has
mK, although the rms noise of our 2 mm spectrum is about 0.5 mK for 4 km s-1 wide velocity channels. Apparently lines
weaker than about 3.5 mK have escaped detection not because of poor
signal-to-noise ratios but because their number is so high that they
are all blended and therefore cannot be easily distinguished.
At mm, 6 lines have been detected from the signal band
in a 256 MHz wide range. Since each line has a width of 9 MHz,
21% of the spectral range observed is covered by detected lines.
The weakest of those have a
of about 1.5 mK, which
corresponds to the 3
noise level at a velocity resolution of 10 km s-1. In contrast to the situation at 2 mm, we have not
yet reached the 3 mm confusion limit. Nevertheless, at 3 mm we
also see signs of weak blended lines at a level of
1 mK.
Weaker lines, of the order of 0.5 mK, have been identified recently
(Cernicharo et al. 2004) making use of the
characteristic line shapes of IRC+10216.
Thus, at the IRAM 30-m telescope the confusion limit toward
IRC+10216 is reached at mm after integrating down to a
system noise of about 1 mK, while at 3 mm one has to integrate
down to 0.3 mK. For other well-known molecular sources, confusion
limits are given in Appendix B.
Transition |
![]() |
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![]() |
mK km s-1 | km s-1 | km s-1 | ||
CS J=2-1 | ||||
C34S | 11 600(100) | -26.2 | 14.0 | 0.55 |
C36S | 96(21) | -26.2b | 14.0b | 0.55b |
C36S/C34S | 8.3(1.8
![]() |
|||
CS J=3-2 | ||||
C34S | 22 190(93) | -26.0 | 14.5 | 0.64 |
C36S | 170(81) | -26.0b | 14.5b | 0.64b |
C36S/C34S | 7.7(3.7
![]() |
|||
SiS J=5-4 | ||||
Si34S | 3180(80) | -26.4 | 14.4 | 0.82 |
Si36S | 41(9) | -26.4b | 14.4b | 0.82b |
Si36S/Si34S | 12.9(
![]() |
|||
SiS J=6-5 | ||||
Si34S | 4280(70) | -26.5 | 14.1 | 0.76 |
Si36Sc | 40(18) | -26.5b | 14.1b | 0.76b |
Si36S/Si34S | 9.3(4.2
![]() |
a The ratio between the intensities at the central velocity of the
line and at the velocity of the horns (i.e. maximum blue or red
shifted emission); b fixed to values obtained for the C34S or
Si34S lines; c simultaneous fit with a blended, unidentified
line at
![]() ![]() |
Ratio | solar | ISM | C.R. | IRC+10216 |
system | ||||
32S/34S | 22.64(0.00) | 24.4(5)b | 16(-5, +11) | 21.8(2.6) |
32S/33S | 126.95(0.05) | ![]() |
38(-18, +500) | 121(15) |
32S/36S | 6519(20) | 3280(760) | 2700(600) | |
34S/33S | 5.606(0.002) | 6.3(1) | 5.6(.3) | |
34S/36S | 288(1) | 115(17) | 107(15) |
a References: solar system data: Canyon Diablo troilite (CDT),
Ding et al. (2001); ISM data: Chin et al.
(1996), Mauersberger et al. (1996), most
ISM data are from the inner Galaxy; Cosmic Rays: these are
"source abundances'', Thayer (1997); IRC+10216: Kahane
et al. (2000), this paper. b The average value for sources at various galactocentric radii; the data by Chin et al. (1996) indicate a possible galactocentric 32S/34S abundance gradient, which would correspond to a value of 32( ![]() |
As already mentioned in Sect. 4 there are emission features at the rest frequencies of both the 3-2 and the 2-1 lines of C36S. These clearly peak above the limits of line confusion. It turns out, however, that at 3 mm there are two emission lines of linear C3H from the image sideband, one of which coincides with the location of the expected C36S 2-1feature. Also in the band is the image of the strong CS 2-1 line. From a well calibrated spectrum of these three lines (Mauersberger et al. 1989) we obtained their relative intensities. According to Cernicharo et al. (2000) there is no hint of variability in time, except for some very highly excited lines or maser lines. In our spectrum of C36S 2-1(Fig. 2) we could therefore apply a fit to the CS line and the two C3H lines from the image band fixing the line shape and the relative intensities. When removing this composite fit from the spectrum, the residual still shows emission at the expected frequency of C36S (Fig. 2). This residual spectrum was used to determine the isotopic ratios given below.
We cannot exclude that the residual line is contaminated since a fit with fixed line shape and rest frequency to the line feature (to the values obtained for C34S) shows some excess emission toward lower velocities. A possible contaminant of the 2-1 line of C36S is a series of K lines of the J=23-22 transition of CH3CCCN (methyl cyanoacetylene; Lovas 1984), which were detected toward a number of Galactic hot cores (Mauersberger et al. 1996). This is, however, unlikely, since this molecule has not yet been identified toward IRC+10216 (e.g. Cernicharo et al. 2000).
We have used the procedure from the CLASS data reduction package to
fit shell type circumstellar lines. In order to obtain more reliable
values for the integrated intensities, we fixed some of the line
parameters of 36S-bearing species to the values obtained from
the more intense 34S-bearing isotopomers. Fit results are given
in Table 1. In the fit to the C36S 3-2 line
feature there is a hint of excess emission toward lower velocities.
At that frequency, the only previously observed interstellar line in
the Lovas (1984) catalog is a transition of CH2CHCN
(vinyl cyanide,
GHz). This complex molecule has
never been detected before in IRC+10216 and the line strength is
very low. It seems plausible that if there is contamination it comes
from a transition of a molecule abundant in late-type carbon stars
but not in the ISM, such as C7H and/or H2C6 (Guélin et al. 1997).
The Si36S 5-4 line can be well fitted, with the determination of the baseline being the major source of uncertainty. The Si36S 6-5 line is also clearly detected, although it seems to be blended with a weak, unidentified feature (Fig. 3). The parameters of the C34S, C36S, Si34S, and Si36S lines and the derived C34S/C36S and Si34S/Si36S intensity ratios are given in Table 1. We have clearly detected C36S and Si36S in IRC+10216. These are the first detections of 36S-bearing molecules in a stellar atmosphere and the first detection of Si36S outside the solar system.
The characteristic line profiles of spherically expanding molecular
shells can be explained in terms of line opacity, expansion velocity
and the relative sizes of the telescope beam and the stellar
envelope (e.g. Kuiper et al. 1976; Olofsson et al.
1982). The slightly U-shaped profiles seen in C34S
and Si34S can be explained by optically thin and spatially
resolved emission. Optically thin lines are expected because the main isotopic lines are only moderately optically thick and because 34S is a factor 20 less abundant than the main isotope
(Kahane et al. 1988; Cernicharo et al. 2000).
Interferometric observations (Guélin et al. 1993; Lucas
et al. 1995, for maps see Grewing 1994)
confirm a compact condensation of about 10'' diameter for main
isotopic CS. SiS is even more compact (Lucas et al. 1995),
which is also apparent by a comparison of C34S and Si34S line shapes (Figs. 2 and 3). It is safe
to assume that C36S and Si36S emission is also optically
thin and it is plausible that the distributions of 34S- and 36S-bearing isotopomers are similar. Chemical fractionation, as
discussed in Chin et al. (1996), should be negligible in
the warm (50 K, Cernicharo et al. 2000) environment of
the inner shell of IRC+10216. Since the isotopomers have very
similar transition frequencies we can assume that the intensity
ratios of the corresponding transitions reflect the abundance ratios
(for a discussion, see Kahane et al. 1988).
From Table 1 it is evident that the relative errors of the integrated line intensities of the 36S-bearing molecules are of the order of 20% or more, while the relative errors for the 34S-bearing species are much smaller. While it is relatively straightforward to determine the X36S/X34S ratios from fitted line profiles (as summarized in Table 1), it is much more difficult to estimate the X34S/X36S ratios since the simple rules for error propagation apply only if the relative errors are small. In addition, if we assume that the error distribution of the X36S/X34S ratios is Gaussian, the distribution of the inverse ratio is not such a simple and symmetric function, making the determination of the variance difficult (see Chap. 3.3.3 in Wall & Jenkins 2003).
The errors as quoted here have been calculated from the formal
errors of the fits to the line profiles. From these four ratios, and
using as weights the inverse squares of the formal errors of the
individual ratios, we have determined the weighted mean of the
36S/34S abundance ratio and its error to be 0.0094(
)
(see e.g. Bevington & Robinson 1992). From a
graphical representation of the individual values and their mean
(see Fig. 5) the only transition whose abundance
ratio deviates from the mean value is the SiS 5-4 line. This might
be an additional hint that for this line there is, in addition to
the statistical uncertainty, also a systematic effect, for example a
blend with an unidentified line (see the discussion above). Since
the relative error of the average 36S/34S abundance ratio
is only 14%, the expectation value of the 34S/36S ratio
and its error can be easily determined. In the following we will use
a 34S/36S ratio of 107(
)
for IRC+10216 to ease a
comparison with values published in the literature.
This value of the best fit is very similar to the interstellar
value of 115()
(Mauersberger et al. 1996)
but is smaller than the solar system value of 288 (Ding et al.
2001). In Table 2, we give a compilation of
all available estimates for sulfur isotope ratios.
36S is one of the more enigmatic stable nuclei
in nature. Partly because of its rarity (the solar system isotopic
abundance is 0.015% relative to 32S), and partly because it is
rarely employed in mass spectrometric studies of sulfur isotopes, up
to very recently even its solar system abundance was quite
uncertain. The often-used Anders & Grevesse
(1989) abundance table lists the S isotopes
with a 34S/36S ratio of 211 (no errors given). An updated
value is recommended in the 1998 compilation of isotopic abundances
of the International Union of Pure and Applied Chemistry (Rosman
& Taylor 1998), where a ratio 34S/
is listed as the "best measurement from a single
terrestrial source''; this composition has also been used by Lodders
(2003) for her Solar System abundance table. Sulfur
isotope ratios measured in the laboratory are commonly reported as
permill deviations from the Canyon Diablo Troilite (V-CDT) standard,
which has an S isotopic composition fairly representative of solar
system materials (IUPAC 2003 compilation; de Laeter et al.
2003). Its composition has recently been precisely
redetermined, with a 34S/36S ratio of
(Ding et al. 2001). Another modification with respect to Anders
& Grevesse (1989) results from a lower
sulfur elemental abundance by a factor 0.86, as first deduced by
Palme & Beer (1993) and quoted in Lodders
(2003) for the solar system. All in all, the solar
abundance of 36S is lower than the commonly-used Anders &
Grevesse (1989) value by a factor of 1.6.
The former uncertainty on this isotope extends to its nucleosynthetic origin. The synthesis of the major isotopes of sulfur, similar to that of other elements between Si and Fe, can mostly be assigned to charged particle reactions during hydrostatic and explosive burning stages in massive stars: 32S and 34S to hydrostatic and explosive oxygen burning, 33S to explosive oxygen and neon burning (Chin et al. 1996; Woosley et al. 2002). For some exceptionally rare neutron-rich isotopes in this mass range including 36S, however, the situation may be different. 36S has a closed neutron shell, hence a small neutron capture cross section, and in an s-process it could, in principle, act as a bottleneck building up a significant abundance relative to neighboring nuclides. For this reason, the weak s-process taking place during core He burning and convective shell C (Ne) burning in massive stars has been suggested as a major source of the 36S abundance in the Solar System (Schatz et al. 1995; Woosley et al. 2002), and according to first network calculations by Schatz et al. (1995) it seemed to be able to fulfill that role. However, re-assessment by Reifarth et al. (2000) with updated nuclear input data indicates a more complex situation. The unexpectedly low neutron capture cross section of 34S causes this nuclide to restrict the flow towards 36S. With rare 35Cl being the main seed nucleus rather than abundant 28Si and 32S, the weak s-process, while actually making 36S, falls short by more than an order of magnitude in quantity in their assessment. Here, a further complication arises because in all recent calculations the abundance of the major seed, 35Cl, which in the s-process zone is given by its original abundance, has been taken from Table 3 of Anders & Grevesse (1989). Actually, there is disagreement for Cl between the nuclide abundances in their Table 3 and the 39% higher elemental abundance given in their Table 1; the latter agrees with other previous and more recent compilations. The predicted 35Cl/37Cl ratio is not affected by these changes, but there is a significant effect on the 36S overproduction factor, which is increased by a factor of 1.4.
According to current models, in massive stars most 36S is produced in a large mass zone where convective shell C-burning operates in hydrostatic conditions before the SN II explosion (Woosley & Weaver 1995). Owing to the high neutron density in the first phase of carbon consumption (Raiteri et al. 1992), more 36S is fed through the neutron channel of the unstable 35S than would be inferred by the occurence of the classical weak s-process. The nucleosynthesis yields of the most recent massive star calculations, made with solar initial composition and with a full network included (Woosley et al. 2002; Rauscher et al. 2002; Limongi & Chieffi 2003), are all based on the Bao et al. (2000) compilation of neutron capture cross sections (where the new and much lower cross section for 34S is included). These results show a production factor of 36S that is typically low by a factor of 2 with respect to the mean of the most abundant nuclides. Taken at face value, this factor of 2 may be reconciled using the new solar 36S abundance as well as the correct 35Cl abundance. However, this does not completely solve the problem of the reproduction of solar 36S, since one would expect the weak s-process products to decline with decreasing metallicity, i.e. earlier generations of massive stars to have been less effective producers of 36S. In addition, there are uncertainties related to the nuclear input data. In particular, the rate of the major channel 36Cl(n, p)36S may be subject to further refinements in the future owing to the much better energy resolution obtained in the new measurements of Wagemans et al. (2003). On the other hand, recent investigations of n-induced reactions on 37Ar and 39Ar (Goeminne et al. 2000, 2002) seem not to affect the 36S production in a significant manner.
A complementary, though probably less important, source of the weak
s-process is the main s-process taking place in the He burning
shell of low mass AGB stars. According to calculations partially
published in Travaglio et al. (2004), the Galactic
contribution to 36S at the epoch of Solar System formation by
all previous generations of AGB stars is about 4%; this may reach 10% once the revised lower 36S and higher 35Cl solar abundances are taken into account.
In Fig. 5 we show the S isotopic data for IRC+10216
(from this work and Kahane et al. 2000) normalized to 34S and solar system ratios. The data clearly show a more than
a factor of 2 enhancement of 36S/34S in IRC+10216 while
the other ratios are solar within the error. An s-process
enhancement of 37Cl in IRC+10216 was found by Kahane et al.
(2000). These authors were able to match the observation
with the predictions of a low mass (1.5
AGB star model
of solar metallicity. The enhancement predicted for the envelope
after the 15th thermal pulse corresponds to the observed Cl isotope
ratio. Re-analysis, using the updated solar abundances, confirms the
agreement, which persists for the last three thermal pulses with
dredge-up (#15 to #17). Included in Fig. 5 are
the corresponding predictions for the sulfur isotopes (assuming, as
in the case of Cl, the updated solar starting composition). It is
obvious in Fig. 5 that qualitatively there is
agreement between model and observations in that in IRC+10216 an
enhancement of only 36S/34S is seen. However, our observed
enhancement of
170% in this ratio is more than a factor of 4 higher than even the new prediction, which already is about twice
as high (
40% instead of
20%) as that in Kahane
et al. (2000).
It has been argued in the past that the central star of IRC+10216
might correspond to an AGB star of intermediate mass, of around 5 ,
instead of the low mass model discussed here. This
possibility was at the upper limit of the range of uncertainty of
the available distance estimates of 110-170 pc (Winters et al.
1994; Crosas & Menten 1997; Le Bertre
1997; Groenewegen et al. 1998; Weigelt
et al. 1998). An intermediate mass was inferred by
Guélin et al. (1995) on the basis of older estimates of
some isotope ratios (see the discussion in Kahane et al.
2000). We have computed a series of AGB models of
different initial mass using an updated network of neutron capture
cross sections and solar abundances. These confirm the conclusions
already reached by Kahane et al. (2000) that an
intermediate mass of 5
is excluded by the almost solar Mg isotope ratios observed. Indeed, in intermediate mass stars the
22Ne(
, n)25Mg and 22Ne(
,
)26Mg reactions would be more
efficiently activated by the higher temperature in the convective He
flashes, providing a clear excess of both neutron-rich Mg isotopes.
Supporting evidence is derived from the observed 37Cl/35Cl ratio, which would be much higher in an intermediate mass star. The resulting 34S/36S ratio, does not depend much on the
initial mass.
While the enhancement of the observed 36S/34S ratio with
respect to the 1.5 models could, in principle, be due to
some unrecognized problems with the AGB models or assumptions on
nuclear parameters involved with it, we believe this to be highly
unlikely and the main s-process in IRC+10216 to be not the main
source of its enhanced 36S/34S ratio relative to solar.
Alternative explanations are briefly discussed below.
![]() |
Figure 5:
Sulfur isotopic ratios in IRC+10216, normalized to
34S and solar ratios. Data from Kahane et al.
(2000) and this work. Also shown are the S isotope
ratios predicted for the envelope of a 1.5 ![]() |
One clue may be the close agreement between the value for IRC+10216
and the one reported for the ISM (also shown in
Fig. 5) by Mauersberger et al.
(1996), which at first glance is puzzling. A
possible solution may be afforded by Galactic chemical evolution
(GCE) and the different times when IRC+10216 and the Solar System
formed, i.e. assuming a starting composition different from solar.
Based mostly on the 12C/13C ratio and limits to C/O,
Kahane et al. (2000) derived an upper limit of 2
for IRC+10216, which nevertheless still allows the
object to be as young as
1.5 Ga. Assuming a correspondingly
evolved starting composition, observations might then be compatible
(within errors) with model predictions. Indeed, first, there is no
need for a large intrinsic contribution to 36S from IRC+10216
(in agreement with the relatively small effect predicted by the AGB star model; Fig. 5). Second, the large difference
between the current ISM (and IRC+10216) and the Solar System is
compatible with massive star model predictions, where the 36S Galactic contribution increases with metallicity (as already pointed
out by Mauersberger et al. 1996). This would
require a rather different evolution of the Cl and S isotopic ratios
during the last few Ga, since the observed 35Cl/37Cl ratio
in IRC+10216 is matched by AGB model predictions using solar as the
starting composition. Indeed, both Cl isotopes are thought to be
synthesized in massive stars in a primary way during explosive Ne and O burning (Weaver & Woosley 1995). Only 4% of the solar
abundance of 37Cl is contributed by the main s-component in low
mass AGB stars, while 35Cl gets actually depleted by neutron
captures, according to the GCE calculations by Travaglio et al.
(2004). On the observational side, 35Cl/37Cl
in Orion has been found to be enhanced relative to solar (Salez et al. 1996), but the errors are too large to reach a firm
conclusion.
A further possibility is that IRC+10216 simply started with a
36S abundance that was unusually high for its time of
formation and location, i.e. invoke inhomogeneity in the ISM. Of
course, it could as well be the Solar System rather than IRC+10216
that started out with an unusual composition. Both suggestions are
not purely ad hoc. Already Lugaro et al. (1999), in
order to explain isotopic variations in the Si isotopes of single
presolar silicon carbide grains from primitive meteorites, called
for small local inhomogeneities in the ISM at the time of parent
star formation of the grains as an alternative to Galactic
chemical evolution. And, as has been known for quite some time
already, oxygen isotopic ratios in the Solar System seem to be
quite peculiar (e.g., Wilson & Rood 1994; Prantzos et al. 1996). In addition, the interstellar Rb isotope
ratio 85Rb/87Rb determined toward
Ophiuchi A
(Federman et al. 2004) indicates a higher ratio of r-
to s-process nuclides in the Solar System as compared to the local
ISM.
Reifarth et al. (2000) discussed various other possible nucleosynthesis scenarios besides the s-process. Accepting Galactic chemical evolution as the reason for the difference between the Solar System and the ISM 34S/36S ratios, those that involve core collapse supernovae (SN) (other than the s-process) face the problem of how to account for the large increase during the last 4.6 Ga. This probably includes the r-process. To some extent this also includes the "neutron burst'', intermediate between s- and r-process and thought to occur during passage of a shock wave through the He shell of type II SN, which Meyer et al. (2000) modelled to explain specific types of isotope abundance anomalies found in supernova grains preserved in meteorites. Note, however, that the efficiency of the process depends on previous s-process seeds and that its effects are included in the SN II yields calculated by the authors cited above. Low-entropy scenarios associated with supernovae type Ia that can account for production of neutron-rich isotopes such as 48Ca and 50Ti (Meyer et al. 1996; Kratz et al. 2001) would be qualitatively compatible with the trend in GCE, but again existing information on production yields for 36S is not encouraging (Reifarth et al. 2000; Kratz et al. 2001).
In summary, while it appears we have observationally identified a stellar source with enhanced 36S and there are some reasonable ideas about how Galactic 36S is produced, we are left with a number of open questions. Further studies of other objects and/or regions in the ISM, especially when correlated with other isotopic systems such as the carbon, oxygen, and chlorine isotopes, may turn out to be helpful. The same holds true for continued study of the relevant stellar physics and nucleosynthesis processes.
Acknowledgements
V. Ahrens from the I. Physikalisches Institut der Universität Köln gave us fits to the line frequencies of C36S prior to publication. R.G. acknowledges support by the Italian MIUR-Cofin2002 Project "Nucleosynthesis in the Early Phases of the Universe''.
Nr. |
![]() |
Freq.a | Im.-freqa | Identification |
mK | MHz | MHz | ||
1 | 4 | 94 903 | 98 126 | HCCC13CN 36-35 |
2 | 3.5 | 94 913 | 98 155 | SiC4 32-31 (im.) |
3 | 5.5 | 95 018 | 98 011 | l-C3Hb (im.) |
4 | 6 | 95 035 | 97 095 | C36S; l-C3Hb (im.) |
5 | 21 | 95 049 | 97 980 | CS(2-1) (im.) |
6 | 20 | 95 087 | 97 942 | SiC4 31-30 |
7 | 1.5 | U95 115 | 97 914 | |
8 | 1.5 | 95 126 | 97 903 | SiCCC13C 32-31 |
9 | 20 | 142 297 | 151 263 | H2C4 161,16-152,15 |
10 | 23 | 142 321 | 151 235 | Al37Cl 10-9c |
11 | 3.5 | 142 345 | 151 215 | HCCC13CN 54-53 |
12 | 5 | 142 402 | 151 158 | SiS J=8-7, v=4d? |
13 | 29 | 142 410 | 151 150 | NaCN 927-826c |
14 | 5 | U142 489 | 151 071 | |
15 | 40 | 142 502 | 151 058 | CCS 1111-1010c |
16 | 7 | 142 525 | 151 035 | C36S 3-2c |
17 | 460 | 142 560 | 151 000 | 29SiS 8-7c |
18 | 4 | U142 594 | 150 966 | |
19 | 12 | U142 622 | 150 940 | |
20 | 7 | 142 675 | 150 885 | SiCCC13C 48-47 |
21 | 9 | 142 695 | 150 865 | artefact |
22 | 850 | 142 730 | 150 831 | C4H N=15-14c |
23 | 9 | 142 755 | 150 805 | C4H N=15-14? |
24 | 850 | 142 768 | 150 793 | C4H N=15-14c |
a The band from which the line probably originates is given in bold face; for information, the frequency of the corresponding image band is also given. b Thaddeus et al. (1985). c Cernicharo et al. (2000) and references therein. d The SiS v=4 J=8-7 line is at 142 399.8 MHz. A shift of 1-2 MHz could be possible if the line is produced in the acceleration region (blue). |
In Table A.1 we list the observed spectral line
features seen in Fig. 1, and the approximate
values for
,
the rest frequency and the rest
frequency in the image band (assuming an LSR velocity of
). The final column contains some identifications
from the literature. In the 256 MHz wide spectrum of the C36S 2-1 line we have identified 8 spectral line features. Of those,
three can be identified with transitions from the image sideband.
In the 256 MHz around the frequency of the C36S 3-2 line we
have detected 16 spectral line features, none of which can be
assigned to a line from the image band. Six lines can be
identified with known transitions from the Cernicharo et al.
(2000) line survey, and the other identifications were
made with the catalog of the Instituto de la Estructura de la
Materia (Madrid) containing
1000 species described by the
authors. There is a clearly detected feature at the frequency of
the C36S line; another feature at 142.755 GHz is tentatively
assigned to emission from C4S. Eight further lines have not
been identified. The emission at 142.695 GHz is an artefact due
to some bad spectrometer channels.
Source | ![]() |
Weakest |
![]() |
Ref.b |
![]() |
detect. lines | |||
mm | mK | h | ||
IRC+10216 | 3 | 1.5 | 28 | 1) |
28 km s-1 | 2 | 3.5 | 10 | 1, 2) |
1.3 | ![]() |
80 | 2) | |
Orion-KL | 3 | 20 | 0.8 | 3) |
5 km s-1 | 2 | 50 | 0.3 | 4) |
1.3 | 70 | 0.4 | 3) | |
SgrB2(N) | 3 | ![]() |
0.4 | 3) |
![]() |
2 | ![]() |
0.4 | 3) |
1.3 | ![]() |
0.5 | 3) |
a Integration time (on+off) with present receivers at the IRAM
30-m telescope under normal winter conditions (good summer
conditions), one polarization to obtain an rms of 1/3
![]() b References: 1) this paper 2) Ziurys et al. (2002); 3) unpublished data; 4) Mauersberger et al. (1988). c Weaker lines can be identified if one makes use of the unique line shapes of the spectra of this source (Cernicharo et al. 2004). |
It is evident from Sect. 5.1 that in our 2 and 3 mm spectra of IRC+10216 we are close to the confusion limit, where an increase of integration time does not yield much further information. It is interesting to investigate where line confusion begins to play a role for other favorites of molecule hunters such as the Orion Hot Core, Sgr B2 or the starburst galaxy NGC 253. We limit our discussion to the IRAM-30 m telescope. If one wants to extrapolate our results to other telescopes with a higher or lower resolution one has to take into account the detailed source structure chemical and physical conditions within the regions observed (Comito & Schilke 2002).
The definition of the useful observing time or rms to be reached
is by no means unique and depends on whether one is interested in
a mere detection of a line or whether one also wants to obtain
some detailed information on the line shape. Here we try to be
pragmatic: we have investigated spectra of several molecular line
sources made with the IRAM 30-m telescope at 3 mm, 2 mm and 1.3 mm wavelength made with such a long integration that
of the spectral range observed is covered with features. We
have determined the antenna temperature of the weakest unequivocal
line features (knowing the typical line width and shape in the
sources studied). The results are given in
Table 5.1. We also give an indication of the
necessary rms for a 3
detection with a velocity
resolution 1/5 the full line widths typical for these sources and
the corresponding necessary integration time with the 30-m
telescope and its present receivers (including all observing
overheads, using one polarization only). In Orion, part of the
confusion arises because of the large line widths observed in the
outflow source. This can e.g. be prevented by observing at a
carefully selected position offset from the outflow source (Combes
et al. 1996).