Abundances in intermediate-mass AGB stars undergoing third dredge-up and
  hot-bottom burning
Mon. Not. R. Astron. Soc. 000, 1–14 (2007) Printed 7 September 2021 (MN LATEX style file v2.2)
Abundances in intermediate-mass AGB stars undergoing
third dredge-up and hot-bottom burning
J.A. McSaveney1,2, P.R. Wood1, M. Scholz3, J.C. Lattanzio2 and K.H. Hinkle4
1Mount Stromlo Observatory, Research School of Astronomy and Astrophysics, Australian National University, Cotter Road,
Weston Creek ACT 2611, Australia. 2Centre for Stellar and Planetary Astrophysics, Department of Mathematical Sciences,
Monash University, Melbourne, Australia. 3 Institut f. Theoretische Astrophysik d. Univ. Heidelberg, Albert-Ueberle-Strasse 2,
69120 Heidelberg, Germany, and School of Physics, University of Sydney, NSW 2600, Australia. 4 National Optical Astronomy
Observatory, PO Box 26732, Tucson AZ 85726-6732, USA.
Accepted 2007 ?. Received 2007 ?; in original form 2007 ?
ABSTRACT
High dispersion near-infrared spectra have been taken of seven highly-evolved,
variable, intermediate-mass (4-6 M⊙) AGB stars in the LMC and SMC in order to
look for C, N and O variations that are expected to arise from third dredge-up and
hot-bottom burning. The pulsation of the objects has been modelled, yielding stellar
masses, and spectral synthesis calculations have been performed in order to derive
abundances from the observed spectra. For two stars, abundances of C, N, O, Na, Al,
Ti, Sc and Fe were derived and compared with the abundances predicted by detailed
AGB models. Both stars show very large N enhancements and C deficiencies. These
results provide the first observational confirmation of the long-predicted production
of primary nitrogen by the combination of third dredge-up and hot-bottom burning
in intermediate-mass AGB stars. It was not possible to derive abundances for the
remaining five stars: three were too cool to model, while another two had strong
shocks in their atmospheres which caused strong emission to fill the line cores and
made abundance determination impossible. The latter occurrence allows us to predict
the pulsation phase interval during which observations should be made if successful
abundance analysis is to be possible.
Key words: stars: AGB and post-AGB – stars: abundances – stars: oscillations.
1 INTRODUCTION
AGB stars are predicted to be major contributors to the
enrichment of the interstellar medium in carbon, nitro-
gen and s-process elements (e.g. Iben and Truran 1978;
Dray et al. 2003). The 12C and s-process elements are
brought to the surface of AGB stars during the third dredge-
up (Iben and Renzini 1983) and they are then ejected into
the interstellar medium by the extensive stellar winds that
occur during the AGB phase of evolution. The 14N produc-
tion is mainly from the envelopes of the more massive AGB
stars where the hot-bottom burning process (Scalo et al.
1975) can cycle the entire envelope through the outer parts
of the hydrogen-burning shell, converting 12C to 14N. In
the case where the 12C in the envelope has a component
that comes from third dredge-up, the 14N resulting from
the dredged-up 12C is of primary origin. There is strong
evidence that some of the 14N in the universe is of pri-
mary origin (van Zee et al. 1998), with the primary source
dominating for metal abundances less than one third so-
lar. Hot-bottom burning in intermediate-mass (M & 3M⊙)
AGB stars is a possible source for this primary nitrogen.
Detailed stellar evolution and nucleosynthesis calculations
suggest that the hot-bottom burning in AGB stars may also
involve the Ne-Na and Mg-Al chains, leading to changes
in the surface abundances of Al amd Mg as well as ni-
trogen (Karakas and Lattanzio 2003; Lattanzio and Wood
2004). Unfortunately, all the evolution and nucleosynthe-
sis calculations are very uncertain due to the problem of
treating convection, especially the critical process of convec-
tive overshoot. Observational constraints from intermediate-
mass AGB stars are required to constrain the convection
parameters such as mixing length and overshoot distance,
which are critical to the nucleosynthesis results. This study
was initiated to provide such constraints.
Considerable numbers of intermediate-mass AGB stars
are known in the Magellanic Clouds (Wood, Bessell & Fox
1983), and these stars can be used to provide the obser-
vational tests required for the evolution and nucleosynthe-
sis calculations. Some work has been done in this area by
Smith et al. (2002) who examined the C, N and O abun-
c© 2007 RAS
http://arxiv.org/abs/0704.1907v1
2 J.A. McSaveney et al.
dances in a sample of luminous AGB stars in the LMC.
They found that 14N was enhanced in a way consistent with
first dredge-up only, without the need for hot-bottom burn-
ing. However, the stars studied by Smith et al. (2002) were
not known variable stars, meaning that they were not near
the end of their AGB evolution where the effects of dredge-
up and hot-bottom burning should be most pronounced. In
addition, the lack of pulsation means that there was no way
to estimate the mass of the stars involved.
In this study, we have aimed to get surface abundances
for a small sample of intermediate-mass, pulsating AGB
stars. These stars are near the end of their AGB evolution
and are about to enter their final superwind stage where
most of the material currently in their envelopes will be
ejected back into the interstellar medium (in fact, two of
the stars are IRAS sources already exhibiting strong super-
winds). Since they are pulsating, we can also estimate their
current masses from pulsation theory. On the other hand,
because of their low effective temperatures and their pulsa-
tion, the model atmospheres required for the interpretation
of their spectra are complicated and difficult, and must in-
volve the dynamics of the atmosphere.
In Section 2, we describe the sample of stars and the
photometry and spectra obtained for them. In Section 3,
we describe the pulsation models, in Section 4 the model
atmospheres are described and the derived abundances are
given. The results are discussed in the final section.
2 OBSERVATIONS
2.1 Near-infrared photometry
The sample of objects observed in this study is given in
Table 1. All these stars are oxygen-rich, pulsating, lumi-
nous AGB stars. All the stars have been monitored in J
and K with the ANU 2.3m telescope at Siding Spring Ob-
servatory, using either a single channel photometer (prior to
March 1994) or the infrared array CASPIR (McGregor et al.
1994).The mean bolometric luminosity Mbol of each star
was derived from the JK photometry and monitoring, using
the bolometric corrections from Houdashelt et al. (2000a,b)
with [Fe/H ] = −0.3 and -0.6 for the LMC and SMC, respec-
tively. Distance modulii of 18.54 and 18.93 and reddenings
E(B-V) of 0.08 and 0.12 were assumed for the LMC and
SMC, respectively (Keller and Wood 2006).
The least evolved of the stars in Table 1, with lowest
luminosity, smallest amplitude and shortest pulsation pe-
riod, is the most luminous red giant in the LMC cluster
NGC1866. At the other extreme are two IRAS sources with
very long pulsation periods and large amplitudes. The lu-
minosities and periods of the stars in the sample mean that
their current masses lie in the range 3–8 M⊙ (Wood et al.
1983).
2.2 The high-resolution near-infrared spectra
Spectra were taken of these stars using the Phoenix spectro-
graph (Hinkle et al. 2003) on the Gemini South telescope.
The observations were done in service mode, and the obser-
vation dates are listed in Table 1. For each star, observations
were made in 3 separate bands, as shown in Figures 1, 2 and
Table 1. The intermediate-mass AGB star sample
Star Mbol P
1 ∆K MC Date2
NGC1866 #4 -6.00 158 0.10 LMC 030210
HV 2576 -6.61 530 0.25 LMC 030209
HV 11303 -5.59 534 0.70 SMC 030728
HV 12149 -6.84 742 0.70 SMC 020920
GM 103 -6.80 1070 1.30 SMC 030728
IRAS 04516-6902 -6.80 1090 1.30 LMC 030211
IRAS 04509-6922 -7.33 1290 1.50 LMC 021202
1 Pulsation period in days.
2 Date of Gemini observation in the form yymmdd. For GM 103
and HV 11303 the observations were spread over 6 days and the
given date is the mean value.
Table 2. The modelled stars
Star L/L⊙ Teff M/M⊙ ℓ/Hp α
NGC1866#4 20000 3490 4.0 2.57 0.05
HV 2576 35000 3350 6.0 2.21 0.51
HV 11303 34000 3490 4.6 2.21 0.58
GM 103 41840 3040 6.0 1.40 1.30
1 Turbulent viscosity parameter.
3. Firstly, an observation centred near 1.554 µm was made
so that the OH lines could be measured for the derivation
of an oxygen abundance; lines of CN also occur in this piece
of spectrum. Secondly, an observation near 2.340 µm was
made to measure 12CO (and possible 13CO) lines and hence
the 12C abundance; given the C abundance, the CN lines
near 1.554 µm then yield a 14N abundance. The Na abun-
dance was also obtained from this spectral region. Thirdly,
an observation at 2.112 µm was made to include lines of Al
(and possible Mg). Lines of Fe, Sc and Ti also occur in the
spectra, giving an estimate of the metal abundance.
3 PULSATION MODELS
Pulsation models and model atmospheres were made for
four stars in Table 2: NGC1866 #4, HV2576, HV11303
and GM103. The spectra of the remaining stars in Table
1 (HV12149, IRAS04516-6902 and IRAS04509-6922) were
considered too difficult to model at this stage because of
the extreme appearance of the spectra, with no clear con-
tinuum remaining (see Figs. 1, 2 and 3). This is presumably
a result of the very large amplitude and very low effective
temperatures of these stars. The J and K photometry for
the four stars for which modelling was attempted is given in
Table 3. Typical photometric errors are less than 0.03 mag-
nitudes. The J and K magnitudes have been converted to
those on the AAO system of Allen and Cragg (1983) using
the conversions in McGregor (1994).
For each observed star, given estimates of the mean lu-
minosity L and the effective temperature Teff from the J and
K photometry, the stellar radius was computed from the
definition L = 4πσR2T 4eff . Then, given the radius and the
known pulsation period, the current stellar mass was com-
puted from the P -M -R relation of stellar pulsation. These
values of mass and luminosity were then used in the nonlin-
ear pulsation calculations for the star.
c© 2007 RAS, MNRAS 000, 1–14
Abundances in intermediate-mass AGB stars 3
1.548 1.55 1.552 1.554 1.556 1.558
CNOH Ti
GM 103
HV11303
HV 2576
NGC 1866 #4
HV 12149
IRAS04509
IRAS04516
Wavelength (µ m)
Figure 1. The spectra of all objects in the region near 1.554 µm. The most important lines used for abundance determination are
marked.
2.104 2.106 2.108 2.11 2.112 2.114 2.116 2.118
Al Fe
GM 103
HV11303
HV 2576
NGC 1866 #4
HV 12149
IRAS04509
IRAS04516
Wavelength (µ m)
Figure 2. The same as Figure 1 but for the region near 2.112 µm.
c© 2007 RAS, MNRAS 000, 1–14
4 J.A. McSaveney et al.
2.332 2.334 2.336 2.338 2.34 2.342 2.344 2.346
GM 103
HV11303
HV 2576
NGC 1866 #4
HV 12149
IRAS04509
IRAS04516
Na Sc CO
Wavelength (µ m)
Figure 3. The same as Figure 1 but for the region near 2.340 µm.
Table 3. JK photometry
JD241 J-K K JD24 J-K K
NGC1866#4 HV11303 (continued)
50731 1.23 9.71 52804 1.29 10.02
50772 1.27 9.69 52924 1.26 9.97
50885 1.24 9.68 53516 1.20 9.71
51157 1.30 9.67 53574 1.20 9.38
52564 1.29 9.71
52714 1.33 9.64 GM103
52803 1.25 9.69 45980 1.66 9.55
52924 1.28 9.69 46341 1.49 8.65
46645 1.54 8.57
HV25762 48142 1.63 9.31
52563 1.39 9.00 48168 1.70 9.23
52715 1.43 9.13 48640 1.64 8.36
52804 1.29 8.97 48851 1.62 8.78
52924 1.36 8.87 48934 1.63 9.13
48992 1.67 9.33
HV113032,3 49259 1.89 9.80
49259 1.19 9.81 49317 1.87 9.82
49291 - 9.61 49374 1.68 9.45
49317 1.17 9.47 52564 1.65 9.33
49374 1.19 9.32 52714 1.59 8.74
52564 1.21 9.27 52803 1.61 8.61
52715 1.39 9.56 52924 1.64 8.40
1 JD24 is Julian Date - 2400000. 2 Extra photometry is given in
Wood et al. (1983). 3 Extra photometry is given in Catchpole and
Feast (1981).
The nonlinear pulsation code described in Keller and
Wood (2006), and references therein, was used in this study.
The turbulent viscosity parameter αν was adjusted to give
the observed K light curve amplitude. The K and V mag-
nitudes were computed using L and Teff from the pulsa-
tion models together with the bolometric corrections from
Houdashelt et al. (2000a,b).
For three of the four modelled stars, MACHO light
curves are available: the V magnitude was computed from
the MACHO magnitude using the transforms in Bessell and
Germany (1999). For the fourth star, HV11303, an Eros2
light curve was kindly provided by Patrick Tisserand (pri-
vate communication), who also provided a conversion from
the two Eros2 bands to the V band. The combined V and
K light curves give a V -K colour curve. The effective tem-
peratures of the models were adjusted (by altering the ra-
tio of mixing-length to pressure scale height ℓ/Hp to repro-
duce the observed V -K colour when available, or the J-K in
the one remaining case of HV11303 (the Eros2 light curve
was obtained after the modelling was completed). The final
model parameters are given in Table 2. In all cases, a helium
abundance Y = 0.30 was assumed, while a metal abundance
Z=0.01 was assumed for the LMC stars and Z=0.004 was
assumed for the SMC stars.
In order to construct a model atmosphere taking
account of the dynamical processes in the outer lay-
ers, procedures similar to those described in Bessell et al.
(1996), Hofmann et al. (1998), Scholz and Wood (2000) and
Ireland et al. (2004a,b) were used, with updates to the
molecular opacity from Ireland and Scholz (2006). Briefly,
c© 2007 RAS, MNRAS 000, 1–14
Abundances in intermediate-mass AGB stars 5
a pulsation model at the same pulsation phase as the phase
of the Gemini observation was extracted from each pulsa-
tion series. The radial profiles of gas pressure and velocity
for this model were transferred to a non-grey, spherical at-
mosphere code, and the temperature structure was iterated
to a new state, now consistent with the atmosphere code.
Once the temperature, gas pressure, density and velocity
structure were all consistent in the atmosphere code, the
model atmosphere structure was transferred to a separate
spectrum synthesis code which was used to compute line
profiles (see Section 4).
In the atmosphere code, half- and quarter-solar metal
abundances from Grevesse et al. (1996) were adopted for
LMC and SMC stars, respectively. Radiative and local ther-
modynamic equilibrium was adopted and, in particular, the
shock-heated region behind the shock front was assumed to
have negligible width and to have no effects on the temper-
ature stratification.
The pulsation series and the extracted dynamical model
used as input to the atmosphere code is now described for
each star.
3.1 NGC1866 #4
NGC1866 #4 is the most luminous red giant in the LMC
cluster NGC1866 (the designation #4 is that given by Fro-
gel, Mould and Blanco 1990: note that Maceroni et al. 2002
have mislabelled #4 as #2 in their paper). The MACHO
light curve of this star shows a period of 158 days. Mod-
elling the cluster HR-diagram, Brocato et al. (2003) derive
masses for the He-burning giants of 3.8–4.15 M⊙. Our de-
rived mass for NGC1866 #4 is 4 M⊙. With this mass and
the computed mean luminosity and Teff , this star must be
pulsating in the first overtone mode, a result also consistent
with the small amplitude.
Fig. 4 shows the observed and computed K, V and Mbol
light curves. The period and amplitude of the pulsation and
the V -K colour are reproduced well. The Gemini observa-
tion was made at a phase of ∼0.1 after visual maximum.
The structure of the pulsation model at this time is shown
in Fig. 5. The pulsation amplitude of this star is quite small
so there are no shock waves present in the atmosphere. The
temperature solution from the model atmosphere program
produces a cooling of the outer layers and a warming at
moderate optical depths compared to the modified Edding-
ton approximation used in the pulsation models.
3.2 HV2576
HV2576 is an LMC star which is more luminous than
NGC1866 #4. It has a much longer period, a larger ampli-
tude, and it pulsates in the fundamental mode. The derived
mass for HV2576 is 6 M⊙.
Fig. 6 shows the observed and computed K, V and Mbol
light curves. Like many of the intermediate-mass, large-
amplitude AGB pulsators in the Magellanic Clouds, this star
shows a double-humped optical light curve. The pulsation
models do not accurately reproduce the double hump but
they do show some evidence for it. The K light curve does
not show evidence for a prominent hump, but the points are
too sparse in the light curve to be definitive. Overall, the
Figure 4. The K, V and Mbol light curves of NGC1866 #4.
The solid points are model values (several cycles are overplotted),
while the crosses are observations or Mbol values computed from
the J and K photometry in Table 3. The V values are computed
from the MACHO MB and MR photometry using the transforms
of Bessell and Germany (1999). The arrows show the phase of
the Gemini observations and the phase of the pulsation model
extracted for model atmosphere computation.
period and amplitude of the pulsation and the large V -K
colour (V -K ∼ 7) are reproduced well. The Gemini observa-
tion was made near minimum light of the K light curve, or
between the double-humped maximum of the optical light
curve. The structure of the pulsation model at this time
is shown in Fig. 7. The velocity gradient through the atmo-
sphere at the time of observation was small and there was no
shock wave present in the atmosphere. As with NGC1866
#4, the temperature solution from the model atmosphere
program produces a cooling of the outer layers and a warm-
ing at moderate optical depths compared to the modified
Eddington approximation used in the pulsation models.
3.3 HV11303
HV11303 is an SMC star which is almost a twin of HV2576
in the LMC in terms of its pulsation period and luminosity.
However, it has a much larger pulsation amplitude. It pul-
sates in the fundamental mode and the derived mass is 4.6
M⊙. This mass is smaller than that of HV2576, which may
explain the larger pulsation amplitude.
Fig. 8 shows the observed and computedK, V and Mbol
light curves. The pulsation models reproduce the observed
K light curve reasonably well. The Gemini observation was
made near minimum light of theK light curve. The structure
of the pulsation model at this time is shown in Fig. 9. Note
that there is an uncertainty of about 0.1 in the phasing of
c© 2007 RAS, MNRAS 000, 1–14
6 J.A. McSaveney et al.
Figure 5. The density ρ, temperature T and velocity v plotted
against radius R in a pulsation model of NGC1866 #4 at the
same phase as the phase of Gemini observation of the star. The
dashed temperature profile is that obtained from the non-grey
model atmosphere code.
Figure 6. The same as Fig. 4 but for HV2576.
Figure 7. The same as Fig. 5 but for HV2576.
Figure 8. The same as Fig. 4 but for HV11303.
c© 2007 RAS, MNRAS 000, 1–14
Abundances in intermediate-mass AGB stars 7
Figure 9. The same as Fig. 5 but for HV11303.
the models relative to the K light curve. With the adopted
phasing, a strong shock was just emerging into the atmo-
sphere at the time of Gemini observation. In addition, there
was a large velocity gradient through the atmosphere above
the shock. The temperature solution from the model atmo-
sphere program shows a large temperature drop in the outer
layers. This behaviour is typical for extended atmospheres
in a certain parameter range where strong water absorption
abruptly dominates in the high layers (Scholz 1985). The
warmer, grey pulsation model used as input to the atmo-
sphere code does not have water-forming low temperatures.
3.4 GM103
GM103 (Groenewegen et al. 1995) is also an SMC star but
of greater luminosity, amplitude and period than HV11303.
It pulsates in the fundamental mode and the derived mass
is 6 M⊙.
Fig. 10 shows the observed and computed K, V and
Mbol light curves. The pulsation models reproduce the K
light curve and the large V -K colour (V -K ∼ 8) reasonably
well. However, the V light curve of the pulsation models
seems to rise towards maximum earlier than the observed
light curve. The Gemini observation was made near maxi-
mum visible light, or about phase 0.1 before maximum of
the K light curve. The structure of the pulsation model at
this time is shown in Fig. 11. This star has a strong shock
situated in the middle part of the atmosphere at the time of
Gemini observation, with moderate velocity gradients above
and below the shock. The temperature solution for GM103
from the model atmosphere program shows a larger devi-
ation from the pulsation solution than in any other star.
This is due to the strong density drop at the shock, with
Figure 10. The same as Fig. 4 but for GM103.
Figure 11. The same as Fig. 5 but for GM103.
c© 2007 RAS, MNRAS 000, 1–14
8 J.A. McSaveney et al.
the model atmosphere’s water-dominated temperature drop-
ping rapidly near what is the effective edge of the star at the
shock front.
4 SPECTRAL SYNTHESIS AND
ABUNDANCE DERIVATION
4.1 Line synthesis calculations
The spherical model atmospheres described in Section 3
were put into a compatible spectrum synthesis program (us-
ing spherical geometry) to generate line spectra. The spec-
trum synthesis program was that used by Scholz (1992),
Bessell et al. (1996) and Scholz and Wood (2000), except
that partial pressures of a larger number of molecules were
calculated using a computer program supplied by K. Ohnaka
(private communication). This program uses an equation of
state based on molecular constants of Tsuji (1973), updated
for CN using data from Costes et al. (1990) and for TiO
from Tsuji (1978). Lines are assumed to be formed in local
thermodynamic equilibrium. As shock-heating behind the
shock front is neglected in these models, emission compo-
nents of synthetic line profiles may only occur as a conse-
quence of large atmospheric extension (P-Cygni-like emis-
sion; see Scholz 1992).
The overall spectrum in each of the three filter re-
gions was computed using a reasonably comprehensive set
of atomic and diatomic molecular lines, although only
specific lines were used in the actual abundance deriva-
tions. For atomic lines, wavelengths, excitation potentials
and log gf values were taken from the Vienna Atomic
Line Database (VALD, Kupka et al. 1999), with param-
eters of lines selected for abundance determination be-
ing checked against those of Smith et al. (2002). Molecular
line data proved more problematic as no single source ex-
ists. CO lines were taken from Goorvitch (1994) and CN
lines from Aoki and Tsuji (1997) and the SCAN database
(Jorgensen and Larsson 1990). CH lines were also obtained
from the SCAN database (Jorgensen et al. 1996). In each
of these cases, lines were cross-checked against those from
Smith et al. (2002). The OH lines proved most problem-
atic. An incomplete selection of lines was assembled from
Smith et al. (2002) and Melendez et al. (2001): several ad-
ditional lines listed in the source Abrams et al. (1994) could
not be not included due to a lack of gf values.
4.2 The derivation of abundances
Starting abundances were adopted by approximating LMC
and SMC metal abundances as half and quarter solar, re-
spectively, with solar defined from Asplund et al. (2005) (see
Table 5). The CO and OH lines computed with these abun-
dances were examined to establish the microturbulent ve-
locity.
As these relatively massive stars are expected to have
altered CNO abundances as a result of nucleosynthetic pro-
cessing, a series of CNO values based on the AGB evolution
calculations of Karakas (2003) and Karakas and Lattanzio
(2003) was adopted, and the resultant spectra were then
calculated. The CNO combination giving the best overall
fit to the observed spectra was used as the primary starting
1.5533 1.5534 1.5535 1.5536 1.5537 1.5538
Wavelength (µ m)
A(O) = 8.13; 8.28; 8.43
OH OH
Figure 12. Fitted OH lines for HV2576. In this and subsequent
figures, fits are shown for the three abundance values shown on
the figure (where A(X) = log[n(X)/n(H)]+12). The dashed lines
correspond to changes of ±0.15 dex from the adopted abundance
(solid line).
1.5533 1.5534 1.5535 1.5536 1.5537 1.5538
Wavelength (µ m)
A(O) = 8.65; 8.80; 8.95
OH OH
Figure 13. Fitted OH lines for NGC1866#4.
point for subsequent refinement of the individual CNO abun-
dances. Once the best CNO abundances had been found, the
metallic abundances were refined to give a best fit. Lines
used for final abundance derivations are listed in Table 4,
and a selection of the fitted lines are shown in Figs. 12 to
The abundances derived from different lines of the same
species show some scatter, with the adopted lines generally
found to yield reasonably consistent values. Figs. 14 and 15
show examples of the fit errors for the CO lines. It is clear
that neither CO line is fitted exactly by the final adopted
abundance (that of the mean of the best fits to 6 lines), but
both lines are fitted reasonably well, with one too strong and
one too weak. Figs. 12 and 13, where Fe I lines are present
on either side of the OH lines, show further examples of
fit errors. The adopted Fe abundance is based on a fit to
the 2.1095µm line in Figs. 24 and 25. However, in HV2576
(Fig. 12), this abundance is too high for the 1.553424µm Fe
line to the blue and too low for the 1.553757µm line to the
red (or, more likely, the models are missing an unidentified
line at this wavelength). For NGC1866#4 (Fig. 13), the
c© 2007 RAS, MNRAS 000, 1–14
Abundances in intermediate-mass AGB stars 9
2.342 2.3422 2.3424 2.3426 2.3428 2.343
Wavelength (µ m)
A(C) = 6.86; 7.01; 7.16
CO CO
Figure 14. Fitted CO lines for HV2576.
2.342 2.3422 2.3424 2.3426 2.3428 2.343
Wavelength (µ m)
A(C) = 7.21; 7.36; 7.51
CO CO
Figure 15. Fitted CO lines for NGC1866#4.
1.553424µm line fits well, while once again the 1.553757µm
line is too weak (or an unidentified line is missing in the
model).
There were also cases where observed and computed
line velocities do not match to better than a few km s−1, as
can be seen in Figs. 14 to 17. The velocity mismatch could
be due to a number of potential causes, including inexact
adopted line wavelengths, inexact wavelength calibration, or
the model velocities as a function of depth not being correct.
The microturbulent velocity required for line broad-
ening varied from star to star. The adopted value for
NGC1866#4 (3 km s−1) is close to the value used by
Smith et al. (2002). However, a larger value of 7 km s−1 was
required for HV2576. This value was derived by fitting to
the CO lines. Slightly different values would be obtained by
fitting different lines. This is probably a reflection of different
amounts of turbulence at different levels in the atmosphere:
the CO lines tend to be broader than the metallic lines and
they are, on average, formed further out in cooler layers.
Our final derived abundances are shown in Table 5. For
the elements Fe, Na, Al, Sc, Ti, and N, the error estimates
in the table are those arising from the fit to the line only, as
only one line was available. They are eye-estimates based on
the fits 0.15 dex above and below the adopted optimum fit
1.556 1.5561 1.5562 1.5563 1.5564 1.5565 1.5566 1.5567
Wavelength (µ m)
A(N) = 8.35; 8.50; 8.65
Figure 16. Fitted CN line for HV2576.
1.556 1.5561 1.5562 1.5563 1.5564 1.5565 1.5566 1.5567
Wavelength (µ m)
A(N) = 8.40; 8,55; 8.70
Figure 17. Fitted CN line for NGC1866#4.
2.34 2.3402 2.3404 2.3406 2.3408 2.341
Wavelength (µ m)
A(Sc) = 2.75; 2.90; 3.05
Figure 18. Fitted Sc I line for HV2576.
c© 2007 RAS, MNRAS 000, 1–14
10 J.A. McSaveney et al.
2.34 2.3402 2.3404 2.3406 2.3408 2.341
Wavelength (µ m)
A(Sc) = 3.25; 3.40; 3.55
Figure 19. Fitted Sc I line for NGC1866#4.
1.554 1.5541 1.5542 1.5543 1.5544 1.5545 1.5546 1.5547 1.5548
Wavelength (µ m)
A(Ti) = 3.45; 3.60; 3.75
Figure 20. Fitted Ti I line for HV2576.
1.554 1.5541 1.5542 1.5543 1.5544 1.5545 1.5546 1.5547 1.5548
Wavelength (µ m)
A(Ti) = 4.55; 4.70; 4.85
Figure 21. Fitted Ti I line for NGC1866#4.
2.3375 2.3376 2.3377 2.3378 2.3379 2.338 2.3381 2.3382 2.3383
Wavelength (µ m)
A(Na) = 5.20; 5.35; 5.50
Figure 22. Fitted Na I line for HV2576.
2.3375 2.3376 2.3377 2.3378 2.3379 2.338 2.3381 2.3382 2.3383
Wavelength (µ m)
A(Na) = 5.60; 5.75; 5.90
Figure 23. Fitted Na I line for NGC1866#4.
2.1088 2.109 2.1092 2.1094 2.1096 2.1098
Wavelength (µ m)
A(Al) = 5.35; 5.50; 5.65
A(Fe) = 7.05; 7.20; 7.35
Al I Fe I
Figure 24. Fitted Fe I and Al lines for HV2576.
c© 2007 RAS, MNRAS 000, 1–14
Abundances in intermediate-mass AGB stars 11
2.1088 2.109 2.1092 2.1094 2.1096 2.1098
Wavelength (µ m)
A(Al) = 5.60; 5.75; 5.90
A(Fe) = 6.90; 7.05; 7.20
Al I Fe I
Figure 25. Fitted Fe I and Al lines for NGC1866#4.
Table 4. Spectral line data.
λ χ Source
(Å) (eV) log(gf)
21095.446 6.20 -0.69 VALD
23378.945 3.75 -0.420 Smith et al. (2002)
23404.756 1.44 -1.278 Smith et al. (2002)
15543.720 1.88 -1.481 Smith et al. (2002)
21093.029 4.09 -0.31 VALD
15535.489 0.51 -5.23 Smith et al. (2002)
15536.707 0.51 -5.23 Melendez et al. (2001)
15560.271 0.30 -5.31 Smith et al. (2002)
15565.815 0.90 -5.00 Melendez et al. (2001)
23396.260 0.37 -5.17 Smith et al. (2002)
23398.224 1.73 -4.44 Smith et al. (2002)
23406.345 0.00 -6.57 Smith et al. (2002)
23408.509 0.36 -5.19 Smith et al. (2002)
23424.328 1.80 -4.427 Smith et al. (2002)
23426.322 0.00 -6.565 Smith et al. (2002)
15563.367 1.15 -1.141 Smith et al. (2002)
abundance. The C and O values are based on averages of the
fits to 6 and 4 lines, respectively, with the stated errors being
based on the standard deviation. The lines used are given in
Table 4. Systematic errors, such as errors in the pulsation
model and its phase, will undoubtedly add significantly to
the fit error, but such additional errors are hard to estimate.
4.3 HV11303 and GM103
As noted in Section 3, the SMC stars HV11303 and GM103
both had strong shock fronts in their atmospheres at the
time of observation with Gemini South. This has made
abundance derivation with our atmosphere code impossible.
Fig. 26 shows a sample of the spectrum for each star, along
with an attempt at synthesisizing the spectrum using the
pulsation models from Section 3.
In the case of GM103, the synthesized absorption lines
are much stronger than the observed lines. A comparison of
the observed and synthesisized spectra shows clearly that
the observed line cores are filled in by strong emission in ev-
ery case. An alternative explanation could be that each line
is split into a pre- and post-shock absorption component,
with no emission in the middle. However, the weakness of
the observed absorption lines compared to the model lines
suggests that there is indeed a strong emission component.
In HV11303, the lines are much broader than in GM103,
with evidence for an emission component as well as absorp-
tion components corresponding to the pre- and post-shock
velocities.
As noted in Section 3, the spectral synthesis code ig-
nores the immediate post-shock, high-temperature region
so that there is no post-shock emission component in our
synthesized spectra. When a strong shock is present in the
atmosphere, this model deficiency prevents abundance de-
termination from lines which have a post-shock emission
contribution.
These results show that observations useful for abun-
dance analysis must be obtained at phases of the pulsation
cycle when no strong shock wave is present in the atmo-
sphere. In these relatively massive stars, the shock wave en-
ters the lower part of the stellar atmosphere near minimum
of the K light curve, or phase ∼0.2–0.3 past minimum vi-
sual light. It has passed through the line forming part of
the atmosphere half a pulsation cycle later, corresponding
to ∼0.1–0.2 in phase past K maximum or after phase ∼0.4
of the visual light curve.
5 DISCUSSION
The abundances obtained here for NGC1866#4 and
HV2576 are shown with other comparable observations in
Fig. 27. The Fe abundance is close to half solar, consistent
with the values obtained by Smith et al. (2002), Hill et al.
(2000) and other studies of young objects in the LMC. The
other metal abundances are similarly near half-solar, except
for Ti in HV2576 where our derived value seems anoma-
lously low.
The CNO abundances are the most interesting part of
this study, since in a highly-evolved, intermediate-mass AGB
star they will be altered by the 1st, 2nd and 3rd dredge-ups
and by hot-bottom burning.
Evidence for hot-bottom burning has been found in
intermediate-mass AGB stars in the Magellanic Clouds by
Plez et al. (1993) and Smith et al. (1995). Plez et al. (1993)
examined spectra of 7 luminous AGB stars in the SMC and
found them to be Li-rich, C-poor and with low 12C/13C
ratios, consistent with hot-bottom burning. Smith et al.
(1995) examined a large number of SMC and LMC stars
and found many luminous AGB stars, including HV2576,
showed strong Li lines that could be attributed to hot-
bottom burning. In fact, HV2576 had the largest estimated
Li abundance of all the stars examined. Maceroni et al.
(2002) examined Li lines in the giant stars in NGC1866
and found that NGC1866#4 (which they mislabel as #2)
c© 2007 RAS, MNRAS 000, 1–14
12 J.A. McSaveney et al.
2.335 2.336 2.337 2.338 2.339 2.34
Wavelength (µ m)
2.335 2.336 2.337 2.338 2.339 2.34
Wavelength (µ m)
Figure 26. Top: A section of the 2.34 µm spectrum of GM103 (points), overlayed with a synthesised spectrum (continuous line). Bottom:
The same piece of spectrum for HV11303.
Table 5. Derived abundances
Star A(Fe) A(12C) A(14N) A(16O) A(Na) A(Sc) A(Ti) A(Al)
HV 2576 7.20 ± 0.10 7.01 ± 0.20 8.50 ± 0.05 8.28 ± 0.08 5.35 ± 0.10 2.90 ± 0.10 3.60 ± 0.05 5.50 ± 0.05
NGC 1866 #4 7.05 ± 0.15 7.36 ± 0.16 8.55 ± 0.15 8.80 ± 0.17 5.75 ± 0.15 3.40 ± 0.15 4.70 ± 0.15 5.75 ± 0.10
Smith et al. 2002 6.37 - 7.16 6.53 - 7.86 7.14 - 8.24 7.82 - 8.33 4.69 - 5.84 2.01 - 2.91 3.84 - 4.61 ...
Half-solar 7.15 8.09 7.48 8.36 5.87 2.75 4.60 5.87
Notes: A(X) = log[n(X)/n(H)] + 12. Half-solar abundances are half the solar abundances in Asplund et al. (2005). The errors are
estimates of the line fitting error alone - see Sect. 4.2.
had enhanced Li compared to other stars in the cluster, in-
dicating that hot-bottom burning is operating in this star
also.
There is evidence that the third dredge-up has been
operating in HV2576, as well as hot-bottom burning.
Wood et al. (1983) give a spectral type M5.5/S1 for
HV2576, indicating the presence of excess ZrO in the spec-
trum, which is a result of the dredge-up of s-process ele-
ments during thermal pulses. Smith et al. (1995) noted that
HV2576 had one of the strongest ZrO bands in their large
sample of stars, indicating that a large amount of third
dredge-up has occurred in this star.
The combination of third dredge-up and hot-bottom
burning is predicted to be a significant source of primary
nitrogen in the universe. The 12C dredged up at a ther-
mal pulse is converted into 14N during the next interpulse
period when hot-bottom burning operates during the quies-
cent H-shell burning phase. Lattanzio and Wood (2004) and
Karakas (2003) (see their Figs. 2.37 and 5.6, respectively)
show the effects of third dredge-up and hot-bottom burn-
ing in intermediate-mass, luminous and sub-solar metallic-
ity AGB stars such as those we are dealing with here. Ini-
tially on the AGB, just at the onset of thermal pulses, hot-
bottom burning by the CN cycle reduces the 12C abundance
to ∼1/15 of its initial main-sequence value, and increases
the 14N abundance to ∼5-6 times its initial value. The ratio
12C/14N at this time is ∼1/15. During subsequent evolution,
the dredge-up of 12C by third dredge-up at thermal pulses,
followed by hot-bottom burning, causes a steady rise over
many pulsation cycles in both the C and N abundance, with
the C/N ratio remaining almost constant at ∼1/15. Dur-
ing this relatively long evolutionary process, there is a small
decrease in the O abundance, but this effect would be too
small for us to reliably detect.
c© 2007 RAS, MNRAS 000, 1–14
Abundances in intermediate-mass AGB stars 13
Looking at Fig. 27 and Table 5, we see that the C abun-
dance in both NGC1866#4 and HV2576 is about 0.1 times
half-solar. At the same time, the N abundance is ∼10 times
the half-solar value. Both these numbers are consistent with
the strong operation of hot-bottom burning in the envelopes
of these stars. Furthermore, the very large N abundances can
only be explained by hot-bottom processing of dredged-up
12C into 14N i.e. a large fraction of the nitrogen in the en-
velopes of these stars must be of primary origin. This is the
first direct demonstration that this long-predicted source of
primary nitrogen does exist.
The O abundance in HV2576 is essentially unchanged
from the half-solar value, as predicted by models. There
is some evidence for an increase in the O abundance in
NGC1866#4. If confirmed, this would lend support to con-
vective theories such as that of Herwig et al. (1997) which
lead to overshoot at convective boundaries, enrichment of
the intershell region in both 12C and 16O at each helium
shell flash, and subsequent dredge-up of both 12C and 16O
in third dredge-up events.
A strong disagreement between the observations and
evolutionary models without convective overshoot relates
to the total stellar mass required for the efficient opera-
tion of hot-bottom burning. The mass of NGC1866#4 must
be no more than ∼4 M⊙ because of cluster membership.
This star shows evidence for efficient hot-bottom burning,
while models without convective overshoot do not predict ef-
ficient hot-bottom burning at masses this low. This suggests
that overshoot (inwards) of convective material at convective
boundaries is required to significantly extends the convec-
tive envelope into nuclear-burning regions. The models of
Ventura et al. (2002) use a convective theory that does pro-
duce overshoot, and their 4 M⊙, half-solar metal abundance
AGB models do have the efficient hot-bottom burning ob-
served in NGC1866#4. At 6 M⊙, both the observations of
HV2576 and models with or without convective overshoot
show efficient hot-bottom burning.
The elements Li, Na and Al can also be compared with
evolutionary models. Smith et al. (1995) give A(Li) = 3.8
for HV2576 while Maceroni et al. (2002) give A(Li) = 1.5
(note that these authors used static model atmospheres, and
observations at random phases where shocks may have been
present in the stellar atmosphere). Models of AGB stars with
efficient hot-bottom burning generally show an initial large
increase in the surface Li abundance, after which it slowly
decreases to low values (A(Li) . 1.0) while the CN cycle
converts C into N - see Karakas (2003) and Ventura et al.
(2002). The high Li value derived for HV2576, together with
the observed high 14N and low 12C abundances, is consistent
with model predictions for a 6 M⊙ star about one quarter of
the way through its TPAGB phase and undergoing efficient
hot-bottom burning - see Fig. C27 of Karakas (2003).
The Na and Al abundances found in this study are
marginally below the expected half-solar values and within
the expected range of observed variation. Models of AGB
stars (Karakas 2003) with efficient hot-bottom burning show
a small increase in Al abundance and small changes in the
total Mg abundance but these are too small to detect here.
The 25Mg/24Mg ratio can increase greatly in such stars, but
we are unable to detect isotopic changes.
0 1 2 3 4 5 6 7 8 9
C N O Fe Na Al Sc Ti
Solar (Asplund et al 2005)
Solar −0.3 (LMC 1st approximation)
HV2576
NGC1866 4
NGC1866 (Hill et al 2000)
Red Giant range (Smith et al 2002)
Figure 27. Final abundances.
6 SUMMARY
High-dispersion, near-infrared spectra have been used to de-
rive abundances of luminous, intermediate-mass AGB stars
in the Magellanic Clouds. The C and N abundances derived
here provide the first observational demonstration that third
dredge-up at thermal pulses followed by hot-bottom burn-
ing can produce significant amounts of primary nitrogen
in intermediate-mass AGB stars. The observed spectra and
modelling show that only spectra taken during the part of
the pulsation cycle near minimum visual light, equivalent to
the latter part of the decline of the K light curve, are useful
for abundance analysis.
The results in this paper come from exploratory obser-
vations and calculations and, in the end, only two stars were
suitable for abundance analysis. It is clear that follow-up
work with a larger sample of objects is required to confirm
the results found here, and to seek C and N abundances as
a function of stellar mass, luminosity and metal abundance.
ACKNOWLEDGMENTS
We would like to thank Amanda Karakas for useful dis-
cussions about AGB nucleosynthesis models, Patrick Tis-
serand for supplying the Eros2 light curve for HV 11303,
and Keiichi Ohnaka for supplying his programs for partial
pressure calculations. JAM is grateful for the discovery grant
from the Australian Research Council (DP0343832) which
supported her during the course of this work: PRW and
JCL also received partial support from this grant. MS re-
ceived support from a grant of the Deutsche Forschungs-
gemeinschaft. This paper is based in part on observations
obtained at the Gemini Observatory, which is operated by
the Association of Universities for Research in Astronomy,
Inc., under a cooperative agreement with the NSF on be-
half of the Gemini partnership: the National Science Foun-
dation (United States), the Particle Physics and Astron-
omy Research Council (United Kingdom), the National Re-
search Council (Canada), CONICYT (Chile), the Australian
Research Council (Australia), CNPq (Brazil), and CONI-
CRT (Argentina). The observations were obtained with the
Phoenix infrared spectrograph, which was developed and
c© 2007 RAS, MNRAS 000, 1–14
14 J.A. McSaveney et al.
is operated by the National Optical Astronomy Observa-
tory. The spectra were obtained as part of Gemini programs
GS-2002A-Q-49, GS-2002B-Q-22 and GS-2003A-Q-25. We
thank the observers at Gemini South for taking these spec-
tra in service mode.
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This paper has been typeset from a TEX/ LATEX file prepared
by the author.
c© 2007 RAS, MNRAS 000, 1–14
http://www.cspa.monash.edu.au/cspa-pub.html
	Introduction
	Observations
	Near-infrared photometry
	The high-resolution near-infrared spectra
	Pulsation models
	NGC1866 #4
	HV2576
	HV11303
	GM103
	Spectral synthesis and abundance derivation
	Line synthesis calculations
	The derivation of abundances
	HV11303 and GM103
	Discussion
	Summary
