A&A 384, 965-981 (2002)
DOI: 10.1051/0004-6361:20020063
J. H. Bieging - M. J. Rieke - G. H. Rieke
Steward Observatory, The University of Arizona, Tucson AZ 85721, USA
Received 4 October 2001 / Accepted 11 January 2002
Abstract
We present spectra covering the wavelength range 2.28 to 2.36 m at
a resolution of
= 0.0007
m (or R = 3500) for a
sample of 24 cool evolved stars. The sample comprises 8 M supergiants,
5 M giants, 3 S stars, 6 carbon stars, and 2 RV Tauri variables. The
wavelengths covered include the main parts of the 12C16O
v = 2-0 and 3-1 overtone bands, as well as the v = 4-2 and 13CO
v = 2-0 bandhead regions. CO lines dominate the spectrum for all the
stars observed, and at this resolution most of the observed features
can be identified with individual CO R- or P-branch lines or blends.
The observed transitions arise from a wide range of energy levels
extending from the ground state to
K. We looked for
correlations between the intensities of various CO absorption line
features and other stellar properties, including
IR colors and mass loss rates. Two useful CO line features
are the v = 2-0 R14 line, and the CO v = 2-0 bandhead. The intensity of
the 2-0 bandhead shows a trend with K-[12] color such that the reddest
stars (K-[12] > 3 mag) exhibit a wide range in 2-0 bandhead depth,
while the least reddened have the deepest 2-0 bandheads, with a small
range of variation from star to star. Gas mass loss rates for both
the AGB stars and the red supergiants in our sample correlate with the
K-[12] color, consistent with other studies. The data imply that stars
with
y-1 exhibit a
much narrower range in the relative strengths of CO 2-0 band features
than stars with higher mass loss rates. The range in observed spectral
properties implies that there are significant differences in
atmospheric structure among the stars in this sample.
Key words: stars: AGB and post-AGB - stars: atmospheres - stars: supergiants - infrared: stars
Cool giants and supergiants typically experience phases of mass loss,
sometimes at very high rates, which may strongly affect the evolution
of the star. In the case of asymptotic giant branch (AGB) stars, the
rate of mass loss exceeds the rate of core growth, so that mass loss is
the most important process controlling the evolution of these stars
from the AGB to the planetary nebula phase. For M-supergiants, mass
loss rates up to 10-4
y-1 are inferred
for timescales of order 105 years,
which implies that the evolutionary consequences may be substantial also
for these massive stars (M > 10
).
The physical mechanisms which drive mass loss are not yet well-understood, though there has been progress very recently with detailed hydrodynamic calculations of models for pulsating AGB stars (Fleischer et al. 1992; Bessell et al. 1996; Winters et al. 1997, 2000; Höfner et al. 1998; Loidl et al. 1999). These models show that the structure of the atmospheres in large-amplitude pulsators is strongly dependent on hydrodynamic processes which include pulsation-driven shocks, non-equilibrium chemistry, and formation of dust grains. Particularly important are the model predictions for time-dependent structure not only over single pulsation cycles but also temporal variations predicted over multiple cycles. The atmospheric structure and its changes with time determine the rate and velocity at which matter is lost, and the amount of dust which is formed and ejected. To understand the mass loss process, then, it is crucially important to compare predictions of the hydrodynamic models with all possible observable properties of such stars to assess the validity of the models.
Comparisons between hydrodynamic model predictions and selected
observational properties have been made in several recent papers.
Winters et al. (1997) successfully predicted the spectral energy
distribution (SED) and light curve at several wavelengths for the
extreme carbon star AFGL 3068, using a self-consistent time-dependent
model incorporating hydrodynamics, chemistry, and dust formation. Hron
et al. (1998) compared low-resolution ISO-SWS 2-15 m spectra of
the carbon star R Scl at two phases with both hydrostatic and dynamic
atmosphere models. They find moderately good agreement with the
dynamic models, and note that the observed time variations require a dynamic treatment of the model atmosphere. Loidl et al. (1999)
calculated synthetic spectra from 0.5 to 12
m for AGB carbon stars
using an opacity sampling method, with the atmospheric structure from
dynamical models. They include 7 major carbon-bearing molecules in the
opacities, and explore the different time dependencies of various
spectral features as a result of the different layers in which
molecules form. They note that their synthetic spectra reproduce
qualitatively the spectra of cool carbon stars, but that higher
spectral resolution is needed in the models to compare with data from
the ISO-SWS, for example.
Bessell et al. (1996) and Scholz & Wood (2000) made hydrodynamical models of M-type Miras and calculated detailed absorption line profiles of CO 1st overtone lines (among others). They compared their model spectra with the velocity results of Hinkle et al. (1982, 1984) to estimate the true pulsational velocities of observed Miras. They find that the CO absorption line profiles exhibit complex variability not only over one pulsational cycle, but also from cycle to cycle. They note the scarcity of observational data at present to compare with their predictions.
Aringer et al. (1999) studied the 4 m first overtone
ro-vibrational bands of SiO in oxygen-rich AGB stars, and compared their
observed spectra with hydrostatic and dynamic atmosphere models. They
confirm the strong variability of the SiO band strengths with light
phase for Miras, and show that this variation is consistent with
pulsationally driven hydrodynamic models. In contrast, hydrostatic
models cannot explain the observed SiO band strengths in cool giants, a
shortcoming which emphasizes the need for dynamical models.
In a high resolution study of molecular hydrogen absorption lines at
2.2
m in 30 Miras and semiregular AGB stars, Hinkle et al.
(2000) find a wide range of line intensities. For a few stars, they
have time-series spectra which show large changes in the H2 1-0
S(1) line intensity over the light cycles, probably reflecting
dissociation and re-formation of H2 with the passage of pulsational
shocks. Such behavior is in fact predicted in the dynamical atmosphere
models of Höfner (1999).
Finally, in the most detailed calculation published to date, Winters et al. (2000) compare velocity-resolved model spectra with high-resolution observations of the CO v = 1-0 fundamental and v = 2-0 first overtone bands in the carbon star IRC+10216. They find that dynamical models for the atmosphere can produce line profiles which are in good agreement with the observations for reasonable model parameters. Even the temporal variations, reflecting the formation and expansion of dust shells, are in rather good agreement with the observed spectra. In other details such as mass loss rate, however, there are discrepancies which suggest that further refinements to the models and more extensive observations are needed.
The CO molecule offers several advantages as a diagnostic probe of the structure of cool stellar atmospheres. It has a stable, closed-shell structure and a high dissociation energy, so it is predicted to form readily in cool giant and supergiant atmospheres, with an abundance close to the chemical thermodynamic equilibrium (TE) value. Even if the gas is being shocked periodically by stellar pulsations, shock chemistry models predict that the CO abundance is scarcely altered from the TE value (Willacy & Cherchneff 1998; Duari et al. 1999). CO should take up almost all available carbon or oxygen, whichever is less abundant, so the CO molecular abundance can be reliably estimated from atomic abundances. The CO ro-vibrational spectrum has bands in the near-IR which are not too seriously affected by telluric absorption. With modern IR detector arrays, the CO bands are accessible to sensitive spectrometers at the fundamental and overtone wavelengths. The molecular constants are well-determined so that line wavelengths and transition strengths can be calculated reliably.
The detailed spectral models for IRC+10216 by Winters et al. (2000)
show that lines of the CO fundamental band at 4.6 m are optically
thick at least in the core, while the first overtone lines, in
particular the v = 2-0 band, are not optically thick, even in models
with large mass loss rates and dense extended atmospheres. Since CO
should form deep in the photosphere and exist essentially unchanged in
abundance out to large distances above the photosphere, this predicted
lack of saturation in the first overtone lines suggests that they
should be very useful probes for the entire extended atmosphere. The
energy level structure of the CO molecule places within a relatively
small spectral window ro-vibrational lines originating from a wide
range in energy levels. Specifically, the intensities of absorption
lines across the v = 2-0 band at 2.3
m probe energies ranging
from the ground state to
K within a spectral range <0.1
m.
For all these reasons, the CO first overtone bands at 2.3 m should
be useful diagnostics of stellar atmospheres models for cool giants and
supergiants. In this paper, we present spectra of selected AGB and
post-AGB stars and late-type supergiants, from 2.28 to 2.36
m.
This wavelength range covers the main part of the v = 2-0 and 3-1 CO
bands, and parts of the v = 4-2 and the 13CO v = 2-0 bands. The
spectral resolution is sufficient to separate a large number of the R-
and P-branch lines, which span a wide range in energy levels. The
spectra should be of value for making comparisons with the predictions of
current hydrodynamic models for cool giant and supergiant atmospheres.
The observations were made with the Steward Observatory 1.54 m telescope on Mt. Bigelow, Arizona (elevation 2510 m) on 1993 October 5 and 6 and 1994 Sep. 12 and 13. The spectrometer was at the Cassegrain focus of the f/45 infrared-optimized chopping secondary.
The spectrometer, described by Williams et al. (1993), has a 120 arcsec
long slit which illuminates a liquid nitrogen cooled
grating. The grating used was ruled at 600 l/mm blazed for
m. The detector was a NICMOS3
HgCdTe array.
The resolution at
m with this grating is
m, or
with a slit width of 1.8 arcsec, which was used for most of the data
presented. During the second observing period, a 3.6 arcsec slit was
used to allow adequate guiding. (As a result the resolution was
degraded to
2500 for a few stars as noted below.) The
grating setting used gave a wavelength range from 2.28 to 2.36
m.
Observations were made by moving the telescope so as to place the star at six positions along the slit. The detector was read out for each position and the separate frames were differenced to subtract the sky and yield two independent spectra of the star. Integration times per exposure ranged from 0.5 s to 30 s, depending on the stellar magnitude.
Standard stars, selected mainly from the Bright Star Catalog, were
dwarfs within a few degrees of the target stars, and with spectral
types between early F and late B. Wallace & Hinkle (1997) have shown
that dwarfs with spectral types from B3 to F5 have essentially
featureless spectra in the wavelength range 2.2-2.4 m when
observed at
3000. Standard and program stars were almost
always observed in pairs at nearly identical airmasses, in an "ABBA"
sequence, i.e., two pairs per observation. Integration time on the
standard stars was typically 30 s per exposure. No absolute photometry
was attempted in either observing period. Wavelength calibration was
determined to first order with exposures of a NeKr lamp, and refined by
an iterative procedure of line identification described below.
|
Spec. | Var. | Obs. | Light | Standard Star | Int. | |||||
Star | Type | Class | Period | K |
![]() |
Datea | Phase | Name | Sp.type | time | RMS |
(d) | (mag) | (![]() |
(cycles) | (s) | |||||||
|
|||||||||||
SU Per | M3.5Iab | SRc | 470 | 1.50 | ... | 236 | -- | HR 870 | F7IV | 24 | 0.037 |
S Per | M3Iae | SRc | -- | 1.31 | 1.4e-6f | 236 | -- | HR 870 | F7IV | 24 | 0.037 |
BI Cyg | M4 | Lc? | -- | 0.58b | ... | 236 | -- | HR 7769 | A2V | 24 | 0.034 |
KY Cyg | M3.5Ia | Lb | -- | 0.28 | ... | 237 | -- | HR 7784 | A1V | 24 | 0.048 |
NML Cyg | M4.5-M7.9 | ... | -- | 0.62 | ... | 237 | -- | HR 8028 | A1V | 24 | 0.033 |
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M2Iae | SRc | -- | -1.88b | 9.1e-8f | 236 | -- | HR 8357 | B6V | 12 | (0.03) |
PZ Cas | M3Ia | SRa | -- | 0.98 | 8.3e-6f | 236 | -- | HR 9019 | A0V | 24 | 0.030 |
TZ Cas | M2Iab | Lc | -- | 1.95b | ... | 236 | -- | HR 9019 | A0V | 24 | 0.037 |
(b) M giants | |||||||||||
W And | M7:p | M | 396.71 | 0.88 | ... | 236 | 0.31 | HR 670 | A1V | 24 | 0.027 |
KU And | M9 | M | 750: | 2.52 | 9.6e-6g | 638 | ... | HR 76 | A0V | 60 | 0.017 |
T Cas | M7e | M | 445.0 | -0.97 | 5.1e-7h | 638 | 0.41 | HR 96 | B9IV | 6 | 0.016 |
IK Tau | M6me | M | 500: | -1.24 | 3.8e-6g | 638 | ... | HR 1137 | A0V | 18 | 0.016 |
UX Cyg | M5 | M | 561.24 | 1.97c | 3.2e-6h | 639 | 0.56 | SAO 70289 | F0V | 36 | 0.012 |
(c) S stars | |||||||||||
R And | S6.6 | M | 408.97 | 0.34 | 1.0e-6g | 236 | 0.00 | HR 63 | A2V | 24 | 0.031 |
W Aql | S4.9 | M | 490.16 | 0.84 | 6.5e-6i | 236 | 0.67 | HR 7366 | A9V | 24 | (0.03) |
AD Cyg | S5.8 | Lb | -- | 1.24 | 2.5e-8i | 237 | -- | HR 7887 | F0V | 24 | 0.035 |
(d) Carbon stars | |||||||||||
HV Cas | CVIIe+ | M | 527 | 2.32 | ... | 236 | 0.1 | HR 343 | A7V | 24 | 0.038 |
V466 Per | N0. V | SR | -- | 1.01 | ... | 236 | -- | HR 1160 | B8V | 24 | 0.025 |
TT Tau | C5II | SRb | 166.5 | 1.05 | ... | 237 | -- | HR 1554 | F2IV | 24 | 0.053 |
V Cyg | C7.4e | M | 421.27 | 0.82 | 2.6e-6h | 236 | 0.34 | HR 7958 | A3V | 24 | 0.032 |
V460 Cyg | C6.3 | SRbj | -- | 0.23 | 1.1e-6h | 237 | -- | HR 8307 | A0V | 24 | 0.039 |
RV Cyg | C6.4II | SRb | 300: | 0.36 | 1.7e-6h | 237 | -- | HR 8307 | A0V | 24 | 0.034 |
(e) RV Tauri variables | |||||||||||
RV Tau | G2 Iaed | RVb | 78.698 | 5.0e | ... | 638 | 0.55 | HR 1554 | F2IV | 360 | 0.007 |
AC Her | F2Ibpd | RVa | 75.4619 | 5.5e | ... | 639 | 0.30 | SAO 103879 | A0V | 360 | 0.009 |
Stars were selected to provide a representative sample of M supergiants and AGB stars of M, S, and C spectral types. Two RV Tauri variables, a class believed to be in a post-AGB evolutionary stage, were also observed. In total, we observed 8 M supergiants, 5 M giants, 3 S stars, 6 carbon stars, and 2 RV Tauri variables. Their properties are summarized in Table 1. Spectral types are from the SIMBAD database, except as noted. Variability class and period are from the General Catalog of Variable Stars (GCVS - Kholopov et al. 1985). K-band magnitudes were from the Two Micron Sky Survey (Neugebauer & Leighton 1969) except as noted. Most of these objects are strongly variable in the visual, but at K-band the variability is much less, typically <1 mag in total range. The gas mass loss rates in column 6 were determined from model fits to observations of the CO J = 1-0 and/or 2-1 mm-wavelength transitions, taken from the literature as noted. For the Miras with well-determined periods, the phase of the light cycle for the observation is given in Col. 8. The period and reference date were taken from GCVS, and we follow the (somewhat inconsistent) convention used there that zero phase is at maximum visual light for the Miras, but is at the deeper minimum in the cycle for the RV Tauri variables. The standard stars, listed in Col. 9, all have spectral types consistent with a nearly featureless spectrum in the observed wavelength range at R = 3500.
The integration time in Col. 11 is the total for all target star exposures averaged into the final spectra. The last column lists an estimate of the rms noise level of the final spectrum. Because almost the entire spectral range covered is filled with absorption features for these stars, it is not possible to calculate a noise level in a region of line-free continuum. Instead, we calculated the ratio spectrum of the difference over the sum of the two observations of a given star in the ABBA sequence, for both the standard and the target, and computed the rms values across these ratio spectra. The rms in Table 1 is then the root-sum-squared of the rms values for standard and target stars, since the final spectrum is the ratio of the target star spectrum divided by the standard spectrum, expressed in terms of the normalized flux. Typical values of the rms noise are in the range 1-3% of the continuum.
The data were reduced with the IRAF software package. Individual exposures (usually 6) were shifted and combined to produce a single two-dimensional spectrum. Flat-fielding was done with a dark current-corrected dome flat. A one-dimensional spectrum was produced with the IRAF aperture extraction tasks, for both the program and standard stars. Finally, the program star spectra were divided by the corresponding standard star spectra to remove telluric features.
Most of the stars were observed at two spectrometer grating settings,
offset by either 0.5 or 1.5 resolution elements. This was done to
ensure that the full resolution was achieved, since the spacing of the
detector elements produced just
sampling of the
spectrum. To combine the spectra taken at two grating settings, the
telluric-corrected spectra of the program stars were interpolated onto
a finer grid in
,
then cross-correlated to determine the exact
shift in pixels between the two grating settings. One spectrum was
shifted with respect to the other by regridding with the offset
determined by cross-correlation, and the two spectra were summed.
Finally, the summed spectrum was multiplied by a normalized 10000 K
blackbody spectrum to correct for the slope introduced in dividing by
the standard star. At
2.3
m, the difference in slope
between a 7500 K and a 15000 K blackbody is negligibly small across
the observed band, so a 10000 K spectrum, appropriate to an A0V star,
was adopted.
An initial estimate of the dispersion functions of the measured
spectra was made from exposures of a neon-krypton lamp. These were not
taken at grating settings identical to all those used for the stars,
however, and the location of the star on the slit introduced a zero-point
offset in the dispersion function, so it was necessary to use an
iterative line identification
procedure to determine an accurate wavelength calibration for each
target star spectrum. As noted in the introduction, the CO
rovibrational bands are potentially useful diagnostic probes of cool
stellar atmospheres because the lines arise from a very wide range of
energy levels but are confined to a relatively small range in
wavelength. The small spectral range means, however, that many lines
are blended or even almost perfectly coincide, especially for wavelengths
longward of the CO 3-1 bandhead. These properties are illustrated in
Fig. 1, which shows the lower state energy of the CO 2-0, 3-1, and
4-2 bands and the 13CO 2-0 band, over the range of wavelengths
covered in the spectra presented here. Wavelengths and energy levels
were calculated from the molecular data of Farrenq et al. (1991).
Shortward of the 3-1 bandhead, the lines are resolved at
m up to about R30. Longward of the 3-1 bandhead,
the lower R- and P-branch lines of the 2-0 band are in some cases
blended and in others just resolved from the R-branch lines of the 3-1
band.
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Figure 1:
Energy levels of CO 1st-overtone bands in the wavelength range
observed. Lower-state energies expressed in temperature units, E/k,
are plotted versus wavelength for the 2-0, 3-1, and 4-2 bands of
12C16O and for the 2-0 band of 13C16O. Selected
transitions are labelled. Horizontal bar at lower left shows
resolution,
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These spectral characteristics of the CO bands led us to an iterative
procedure to refine the wavelength calibration and to identify as many
spectral features as possible. First, the 2-0 and 3-1 bandheads are
usually readily identified. The 3-1 head, however, is blended with
the 2-0 R12 line, but the R14 line lies shortward of the 3-1 head
just enough to be well resolved at R = 3500. Higher R-branch lines up
to about R31 are well-separated, so we used the 2-0 R14 and R26 lines,
and the 2-0 bandhead wavelengths. On the longward side, the 2-0 R0
line, blended with the 13CO 2-0 bandhead was always recognizable
and usually the 4-2 bandhead (but blended with 3-1 R12). These
identifications were used to derive a dispersion function from a
3-segment linear spline fit with IRAF. Because the bandheads are
blended and have shapes which should depend on the properties of the
stellar atmospheres, we made a second iteration using only lines that
were unblended at
m, or that were
coincident with another line. The second pass lines included 2-0 band
R6, R7, R8, and R14 through R31; 3-1 band R15, R16, and R21; 2-0 P4
blended with 3-1 R10; and the 2-0 bandhead. Typical scatter about
the fitted dispersion function was <
m peak-to-peak or an
rms scatter of <
of a resolution element.
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Figure 2: Spectrum of SU Per (spectral type M3.5Iab, variability class SRc). CO lines up to R40 are identified at the bottom, with bandheads as thick vertical lines. Atomic lines from Hinkle et al. (1995) are indicated at the top. |
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Figure 3: Spectrum of S Per (M3Iab, SRc). |
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Some of the atomic lines which Hinkle et al. (1995) detected in their high resolution spectrum of Arcturus are also evident in our spectra. These lines are heavily blended with CO lines except at the edge of the spectra shortward of the 2-0 bandhead, where the noise level increases, so no atomic lines were used in the fits to determine the dispersion functions.
The final spectra are shown in Figs. 2-25,
grouped by stellar category. The flux scale is normalized to unity at
the peak at
2.293
m, just shortward of the v = 2-0
bandhead. The estimated 1-
noise level is shown by a vertical
error bar on the left side. The resolution is
0.0007
m, or R = 3500. At this resolution, the spectra exhibit a
rich complexity of absorption features which are predominantly due to
the v = 2-0 and 3-1 rovibrational lines of 12C16O, over the
wavelength range covered. To aid in identifying individual features in
these CO bands, thin vertical lines mark the wavelengths of the P1-P4
and R0-R40 lines of the v = 2-0 band, and R10-R40 of the v = 3-1
band. Also shown as thicker vertical lines are the heads of the
v = 2-0, 3-1, and 4-2 bands, and the 13C16O v = 2-0 band.
CO line wavelengths were calculated from the molecular constants in
Farrenq et al. (1991). Atomic transitions which were detected in the
high-resolution spectrum of Arcturus by Hinkle et al. (1995) are marked
with lines labelled above the spectrum.
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Figure 10: Spectrum of W And (M7:p, Mira). |
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Despite the apparent complexity in Figs. 2-25, close examination of
each spectrum shows that almost every feature can be identified with
either a single CO line or a blend of lines. The lines from R14 to
about R35 in the 2-0 band are clearly resolved, and the R13 line is
apparent as an inflection just shortward of the v = 3-1 bandhead.
Longward of the 3-1 bandhead, the R-branch lines of the 2-0 and 3-1
bands are alternately blended and resolved. For example, the 2-0 R9
line blends with the 3-1 R33 and R34 lines, and 2-0 R7 with 3-1 R20,
producing at this resolution characteristically deep, broadened
absorption features. In contrast, the gap between the 2-0 R0 and P1
lines allows for a clear separation of the 3-1 R15 and R16 lines. It
is quite possible that some weak absorption is present from lines at
energies higher than the bandheads, i.e., >R51. Hinkle et al.
(1995) detected up to the 2-0 R99 line in the spectrum of Arcturus,
which however has an effective temperature of 4320 K, significantly
higher than the stars in our sample. The higher energy lines may
contribute weakly but the dominant CO absorption features are those
identified in the figures.
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Figure 15: Spectrum of R And (S6.6, Mira). |
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Figure 18: Spectrum of HV Cas (CVIIe+, Mira). |
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Figure 22: Spectrum of V460 Cyg (C6.3, SRb). |
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Figure 25: Spectrum of AC Her (F2Ibp, RV Tauri variable). |
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Only a few atomic lines are readily discernible in the spectra. The
most obvious are a line of Ti I at 2.29696
m and a line
of Sc I at
2.29926
m, both of which show up as absorption
features on the flank of the v = 2-0 bandhead. A few weak atomic lines
are present at the short wavelength edge of the spectra, but other
atomic lines, if present, are badly blended with CO lines.
v = 2-0 | v = 3-1 | v = 4-2 | ||||||||
BH | R30a | R14 | R0b | BHc | R16 | R15 | BHd | |||
Star | ||||||||||
(a) Supergiants | ||||||||||
SU Per | 35 | 68 | 64 | 42 | 32 | 56 | 59 | 35 | ||
S Per | 57 | 67 | 61 | 54 | 45 | 62 | 66 | 53 | ||
BI Cyg | 37 | 73 | 69 | 46 | 37 | 60 | 64 | 41 | ||
KY Cyg | 40 | 65 | 61 | 48 | 36 | 56 | 60 | 42 | ||
NML Cyg | 66 | 79 | 73 | 73 | 60 | 80 | 79 | 70 | ||
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44 | 76 | 76 | 54 | 42 | 62 | 69 | 47 | ||
PZ Cas | 41 | 66 | 61 | 49 | 37 | 56 | 56 | 40 | ||
TZ Cas | 38 | 70 | 68 | 41 | 32 | 56 | 60 | 33 | ||
(b) M giants | ||||||||||
W And | 30 | 77 | 64 | 38 | 33 | 58 | 61 | 35 | ||
KU And | 65 | 81 | 81 | 60 | 62 | 72 | 76 | 77 | ||
T Cas | 45 | 81 | 74 | 49 | 40 | 66 | 71 | 43 | ||
IK Tau | 68 | 91 | 80 | 65 | 64 | 85 | 89 | 64 | ||
UX Cyg | 64 | 85 | 78 | 62 | 56 | 75 | 80 | 59 | ||
(c) S stars | ||||||||||
R And | 31 | 72 | 69 | 34 | 31 | 53 | 59 | 36 | ||
W Aql | 35 | 68 | 66 | 39 | 35 | 52 | 53 | 38 | ||
AD Cyg | 27 | 68 | 66 | 40 | 22 | 57 | 56 | 27 | ||
(d) Carbon stars | ||||||||||
HV Cas | 79 | 84 | 84 | 86 | 80 | 89 | 86 | 82 | ||
V466 Per | 50 | 70 | 74 | 56 | 48 | 65 | 63 | 51 | ||
TT Tau | 34 | 65 | 67 | 46 | 35 | 61 | 63 | 41 | ||
V Cyg | 70 | 87 | 84 | 83 | 74 | 90 | 89 | 79 | ||
V460 Cyg | 39 | 70 | 73 | 48 | 38 | 59 | 59 | 39 | ||
RV Cyg | 52 | 73 | 73 | 61 | 50 | 70 | 66 | 57 | ||
(e) RV Tauri variables | ||||||||||
RV Tau | 80 | 89 | 92 | 97 | 80 | 99 | 100 | 86 | ||
AC Her | 93 | 96 | 95 | 103 | 93 | 102 | 100 | 95 |
The spectra in Figs. 2-25 show broadly similar structure in the CO bands, but there are significant differences from star to star, in the shape and depth of many features. To facilitate comparisons and search for correlations with stellar properties, we list in Table 2 the intensities of selected features, expressed in per cent of the continuum, as measured directly from the spectra. Features tabulated include the following:
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Figure 26:
Comparison of spectrum of ![]() |
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The most directly comparable published spectra for any stars in our
sample are those of Kleinmann & Hall (1986) and Wallace & Hinkle
(1997), who present spectra for 4 stars in common:
Cep, SU Per,
KY Cyg, and PZ Cas - all M supergiants. (Kleinmann & Hall included
Cep and SU Per in their atlas and in fact the same observational
data were processed by Wallace & Hinkle 1997.) We have compared our
spectra with those of Wallace & Hinkle (1997) and find good agreement,
considering the difference in effective resolution. Wallace & Hinkle
used the Fourier Transform Spectrometer (FTS) on the NOAO 4-m telescope
and presented spectra with an effective frequency resolution of 1.6
cm-1 after processing, corresponding to R = 2700 or
m at
m. At this somewhat
lower resolution than ours (R = 3500), the CO bandheads show virtually
identical intensities and shapes but the individual rovibrational lines
(e.g., 2-0 R14-R30) are less well-resolved, as expected. An example
which compares spectra for
Cep is shown in Fig. 26. These
observations were separated in time by over 12 years, so spectral
variability is a possibility, but the main differences can be explained
simply by the difference in effective resolution.
Kleinmann & Hall (1986) noted that the CO 2-0 and 3-1 bandhead
intensities were highly correlated for the stars in their atlas. We
find the same correlation in our data, as can be seen in Table 2.
The 3-1 bandhead at our resolution is somewhat
blended with the 2-0 R12 and R13 lines, which must affect the measured
intensity of the 3-1 bandhead. The 2-0 and 3-1 bandheads have nearly
equal absorption depths in all our stars, spanning a range from 27% to
93% of the continuum. The mean ratio of
I(3-1 BH)/I(2-0 BH) for all
stars in Table 2 is
,
where I(F) is the normalized
intensity of feature F from Table 2. There is no significant
difference between stellar categories. Individual transitions in the
v = 2-0 band and shortward of the 3-1 head also show strong
correlations - e.g., the R14 and R30 line intensities have a mean ratio
I(R30)/I(R14) =
,
with no difference between
categories. (Note that the R30 line coincides with the R71 line.) The
intensity of the 2-0 R0 feature is less well-correlated with that of
R14, but the R0 line is blended with the 3-1 R17 line and the
13CO 2-0 bandhead, so a larger scatter is not surprising.
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Figure 27: v = 2-0 R14 intensity vs. 2-0 bandhead intensity. Symbols denote red supergiants, M giants, S stars, carbon stars, and RV Tauri variables as indicated. Circles indicate Mira variables. |
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A more interesting relation is found in comparing the intensity of the
2-0 R14 line with the 2-0 bandhead, as shown in Fig. 27. There
appears to be a kind of saturation effect in that as the 2-0 bandhead
gets deeper (
), the R14 line approaches a limiting
value of
60% of the continuum. The trend in Fig. 27 seems
well-defined for the AGB stars, but the red supergiants generally lie
below or close to the AGB stars. The 9 Mira variables in our sample
are indicated by circles in Figs. 27-33. In Fig. 27, the Miras show
a tight linear correlation, but the non-Mira AGB stars follow the same
trend with no more scatter than the Miras. For the plotted quantities,
then, there seems to be no difference between variability classes for
the AGB stars.
![]() |
Figure 28: Ratio of R14 absorption depth to (2-0) bandhead absorption depth, vs. (2-0) bandhead intensity. Symbols as in Fig. 27. |
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An alternative way to compare
these features is to plot the ratio of the absorption depths vs.
I(2-0 BH), where the absorption depth (in per cent) is defined as
100 - I(F). Figure 28 shows the ratio of the absorption depth of
R14 over the 2-0 bandhead. For the stars with the deepest 2-0
bandheads (), the depths of the R14 line all cluster near 0.5
of the depth of the 2-0 head. The ratio appears to show an increasing
scatter for shallower 2-0 bandheads (>40%). Taking the AGB stars
and the red supergiants separately, however, Fig. 28 suggests that
there are two different correlations present-one for the giants and
another for the red supergiants. The Miras and non-Mira AGB stars in
Fig. 28 are not distinguishable in the average trend nor the scatter
about it. The two RV Tauri variables do not obviously fit with either
trend.
The infrared color K-[12] should be a measure of the dust opacity of
the circumstellar envelope for cool mass-losing stars. Whitelock et al. (1994) inferred a relation between
and K-[12]
colors for AGB stars; Josselin et al. (2000) found a similar relation
for M supergiants. LeBertre (1997) and LeBertre & Winters (1998)
derived relations between the gas mass loss rates and IR colors for
carbon stars and oxygen-rich (i.e., M-type) Miras, respectively. Their
mass loss rates were determined from radiative transfer models for the
1-100
m spectra, and assumptions about the dust properties.
![]() |
Figure 29:
Gas mass loss rate from CO mm-wavelength emission lines (see Table 1), vs. K-[12] color. Symbols as in Fig. 27. Dotted line shows mean relation found by Whitelock et al. (1994) for Miras in South Galactic Cap. Dashed line is for red supergiants assuming a gas-to-dust ratio of 200 and the
![]() |
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We have calculated the K-[12] colors for our sample using the IRAS PSC
12 m fluxes, S12, and [12] = -2.5 log(
S12/28.3 Jy)
(see the IRAS Explanatory Supplement - Beichman et al. 1985). In Fig. 29, we show the gas mass loss rates for the 13 AGB stars and
supergiants in our sample with published
values (see
Table 1), plotted against K-[12] color. There is indeed a fairly good
correlation as would be expected if the dust mass loss rates determine
K-[12] and the gas-to-dust ratios are similar for these stars. The
scatter at a given K-[12] color does suggest a factor
5 spread
in
,
which is comparable to the spread in gas-to-dust
ratios inferred for samples both of supergiants (Josselin et al. 2000)
and of AGB stars (Whitelock et al. 1994). If we consider only the 7
Miras for which CO-derived gas mass loss rates are available (circled
symbols in Fig. 29), the data points all lie on a straight line, but
with a shallower slope than the relation of Whitelock et al. (1994).
The trend for all the Miras in Fig. 29 is actually nearly parallel to the
curve of LeBertre & Winters (1998) for carbon stars, even though
4 of the 7 points are M-giants and 2 are S stars. Since there is only
1 carbon star in this small sample, we cannot draw any conclusions
about differences between M-giants and carbon stars in the
-
(K-[12]) relation. Figure 29 does suggest, however, that the M-giants
follow a trend which is significantly shallower than predicted by
LeBertre & Winters (1998) for their O-rich dust models.
![]() |
Figure 30: (2-0) bandhead intensity vs. K-[12] color. Symbols as in Fig. 27. |
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If the K-[12] colors are indicators of the gas mass loss rate, then the CO band features might be expected to correlate with the infrared colors as well. In Fig. 30, we plot the intensity of the (2-0) bandhead against K-[12] color for all the stars in our sample. Clearly there is no tight correlation between the two quantities. There is, however, a trend in the data that the observed range in (2-0) bandhead intensities depends on the K-[12] color. At the reddest colors, >3 mag, the (2-0) bandhead ranges from 30% to >80% of the continuum. For less reddened stars (K-[12] <3 mag), the observed range is reduced to 27% to 52%. Conversely, the deepest (2-0) bandhead values (<40% of the continuuum) correspond to the widest range in K-[12] colors, from 0.2 to 5 mag, while the shallowest (2-0) bandheads (>60%) have colors in the range 3-6 mag. If the K-[12] color is indeed an indicator of mass loss rate for AGB and RSG stars (but not the RV Tauri stars), as Fig. 29 suggests, then the distribution of points in Fig. 30 implies that stars with the higher mass loss rates exhibit a wider range of CO bandhead absorption than stars with lower mass loss rates. Lower rates correspond to the deepest CO absorption, with (2-0) bandheads 25%-50% of the continuum.
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Figure 31: Ratio of R14 absorption depth to (2-0) bandhead absorption depth), vs. K-[12] color. Symbols as in Fig. 27. |
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A similar trend is found if we compare the ratio of the 2-0 R14 and
2-0 bandhead absorption depths as a function of K-[12] color. Figure 31
shows that for stars with K-[12] < 3 mag, the R14 line depth is
between 0.43 and 0.57 of the 2-0 bandhead depth. For stars with
K-[12] > 3 mag, the ratio spans 0.45 to 0.92, i.e., about a 3 times
larger range. If we interpret the K-[12] colors in terms of mass loss
rates, 3 mag corresponds to
y-1. The data imply that stars with mass loss rates
lower than this value show a much narrower range in relative strengths
of the CO features than do stars with
y-1.
![]() |
Figure 32: Ratio of R14/(2-0) bandhead intensities vs. K-[12] color. Symbols as in Fig. 27. |
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An alternative comparison of the CO features with IR color is to plot the ratio of the intensity of the 2-0 R14 line to the intensity of the 2-0 bandhead, vs. K-[12], shown in Fig. 32. An important question in evaluating the spectra is whether the CO absorption features are being filled in by a dust continuum. If this effect were significant, the intensity ratios of features such as [2-0 R14/2-0 BH] should approach unity as the degree of "veiling'' increases. Figure 32 shows a marginal trend in this direction, in that for K-[12] > 3 mag, the mean ratio is lower than for stars with K-[12] < 3 mag. The correlation is weak, however, and at any given K-[12] color, there is a significant spread in the [R14/2-0 BH] intensity ratios. This spread in the ratios suggests that dust continuum is probably not a major factor in determining the relative intensities of the CO absorption features.
Figure 33 makes a direct comparison between the 2-0 bandhead and the
CO-derived gas mass loss rates for the 13 stars with
in
Table 1. The trends are very similar to those in Fig. 30, i.e., that
stars with the deepest CO bandheads (
)
span a wide range in
mass loss rates, while the shallower bandheads (
)
are limited
to the highest mass loss rates. This similarity in the trends between
Figs. 30 and 33 is to be expected, given the rather good correlation
between K-[12] and
seen in Fig. 29.
The Mira variables in Figs. 30-33 (circled symbols) show the trend
noted above. For the reddest stars (K-[12] > 3 mag), the ranges in CO
line intensities and ratios are very large, while the less reddened
Miras have more restricted values of CO line properties. Since the
K-[12] color correlates well with
for the Miras (see Fig. 29),
this result implies that Miras with
(5-10)
y-1 exhibit a larger range in atmospheric structure
than do the Miras with lower
.
This dichotomy may possibly be
related to effects found by Winters et al. (2000) in their dust-driven
hydrodynamic models. At high mass loss rates (>
y-1), they find a large range in variation of the
envelope structure with time - e.g., cycle-to-cycle and multiperiodic
behavior. In contrast, for lower mass loss rates, the envelopes are
nearly stationary with time, so that more uniform atmospheric
absorption line properties would be expected. In this theoretical
context, then, the large range of CO properties at high
or
large K-[12] color (Figs. 30-33) is in fact a symptom of the large
variation in atmospheric dynamical structure with time. This
connection reinforces the case for the CO bands as diagnostics for
hydrodynamic models.
![]() |
Figure 33: (2-0) bandhead intensity vs. gas mass loss rate. Symbols as in Fig. 27. |
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There is previous evidence for variability in CO and other molecular
bands in the near IR spectra of cool giants which, it has been thought,
were modulated with the phase of the light cycle, though the connection
with phase is more an assumption than a demonstrated correlation.
Frogel (1971) conducted an extensive study of 18 Mira variables,
mostly of late-M spectral type. His data typically extend over one
period of stellar variation. Changes in CO absorption strength were
detected for most of the stars. The bands were systematically weaker
at stellar temperature minima than at maxima, but there was a wide
range of behavior relative to precise phase of the light cycle. Time
variations in the CO 2-0 bandhead of the S star Cyg were
reported by Wallace & Hinkle (1997), who found that the absorption
depth varied over the light cycle from 0.49 to 0.66 of the continuum
when observed at R = 2700. Aringer et al. (1999) found that for a
sample of 8 Miras, the equivalent width of the SiO first overtone
bandheads varied systematically with phase.
We plotted the CO 2-0 bandhead intensity versus phase for the 9 Miras in our sample with allegedly accurate periods, where the phase was determined from periods and reference epochs in the GCVS. There is no evidence of a correlation in our data. Three factors could mask any modulation of the CO features with the light cycle however. First, there may be a large intrinsic spread in the CO 2-0 bandhead depth from star to star, both in the mean and the extreme values over a cycle. Second, the phases we derived (see Table 1) may have some accumulated error over the time interval from the reference epoch, or the period may have changed or be modulated by a second period. Third, long-term cycle-to-cycle variations in the atmospheric structure may also modulate the CO absorption lines. Recent theoretical calculations find that the envelope structure for some models may change over a timescale longer than the stellar pulsation. Hofmann et al. (1998) see such cycle-to-cycle behavior in their pulsational models for M-type Miras. Winters et al. (2000) find that hydrodynamic models for the carbon star IRC+10216 show individual CO lines varying in shape and intensity over several light cycles, as a consequence of episodes of dust shell formation and ejection. These models suggest that there may be no simple correlation in spectral variations with the light cycle, at least for the Miras. Rather, longer term variations are to be expected, as the observations of IRC+10216 presented by Winters et al. seem to indicate.
Comparison of the spectra of the 4 supergiants in our sample which were
also observed by Wallace & Hinkle (1997) shows no evidence for
variability. The 2-0 and 3-1 bandheads are resolved in both data
sets, and in all 4 stars, the depth and width of the bandheads are the
same within the noise for each pair of spectra. The individual
R-branch lines differ only to the degree expected due to the different
reolutions of the two sets of spectra. The Wallace & Hinkle (1997)
spectra were obtained in 1981 June (
Cep and SU Per) and 1984
April (KY Cyg and PZ Cas), so on a timescale of more than a decade, we
see no evidence for spectral variability in the CO bands for these red
supergiants.
The two RV Tauri variables, which are believed to be in a post-AGB
evolutionary stage, are more interesting candidates for spectral
variability. Oudmaijer et al. (1995) note that 3 of 5 known post-AGB
stars with 2.3 m CO in emission show spectral variability. One of
these is AC Her (= HD 170756). Oudmaijer et al. argue that the CO
overtone bands are in emission at the optical (V-band) minimum, and go
into absorption during the decline after secondary maximum. Their
data clearly show dramatic changes in the CO first overtone spectra,
but the sampling of the light cycle is too sparse and irregular to draw
firm conclusions. In our spectrum of AC Her (Fig. 25), which was taken
at about phase 0.3 near or just after the secondary minimum in the
light cycle, the CO 2-1 bandhead depth is about 7% of the continuum
(i.e., I(2-0 BH) = 0.93). This value is only half the absorption
depth measured by Oudmaijer et al. (1995) for a spectrum taken at phase
0.90. Figure 25 also shows an indication of weak emission in the
vicinity of the 2-0 R0 and P1 lines, at a few per cent of the
continuum, but the bandheads are definitely in absorption. Our
spectrum evidently shows the CO bands in a state intermediate between
the deep CO absorption and emission reported by Oudmaijer et al. They
argued that the variability was a result of episodic mass ejection
during the pulsational cycle, with emission occurring just after
ejection of gas while it is still close to the star in a dense warm
state. Post-AGB mass loss is probably closely related to the formation
and shaping of planetary nebulae (e.g., Kwok 1993), but the physical
mechanisms are poorly understood. Well-sampled monitoring of the 2.3
m CO bands over the light cycle of AC Her and related stars could
be a very useful diagnostic for the mass loss. The relatively short
periods (
)
of these stars would help make such a monitoring
program practical.
The 2-0 bandhead of 13CO at 2.3448 m is a relatively
prominent feature in most of the spectra in Figs. 2-25. It is
important to note, however, that at a resolution of R = 3500 the
bandhead is affected by other lines close in wavelength. The 2-0 R0
and 3-1 R17 lines of the main CO isotopomer, which are almost
coincident, lie within one resolution element of the 13CO 2-0
bandhead. The individual R-branch lines of the main isotopomer are
typically quite strong, and these two clearly contribute to the
absorption at the wavelength of the 13CO bandhead. Interpreting
the 13CO 2-0 bandhead strength in terms of a 12CO/13CO
abundance ratio must include a calculation of the main
isotopomer R-branch lines near the bandhead.
In this connection, there is also a very close coincidence between the
13CO 2-0 bandhead and a line of Ti I at 2.34479 m.
This line is not separated from the 13CO bandhead even at very
high resolution in the spectrum of Arcturus by Hinkle et al. (1995).
Another Ti I line at 2.2970
m is prominent on the side of
the CO 2-0 bandhead in most of our spectra, typically absorbing
5% of the continuum at R = 3500. The strength of this line
suggests that the Ti I line coincident with the 13CO 2-0
bandhead also contributes significantly to the observed absorption
feature. As with the R-branch lines discussed above, the effect of
Ti I absorption must be included in comparing the 13CO
2-0 bandhead with 12CO.
The spectra of our sample of carbon stars (Figs. 18-22) show
remarkable variations in the depth and shapes of the CO 2-0 bandheads
and main R-branch lines (compare for example HV Cas - Fig. 18 - with
V460 Cyg - Fig. 22). The main features in all cases do appear to be
the CO overtone band lines, but it is clear that the bands are
modulated by other molecules with features in the observed range of
wavelengths. It is well known that carbon stars have a rich organic
chemistry in their atmospheres and circumstellar envelopes. A recent
study of ISO-SWS spectra by Jørgensen et al. (2000)
identified C2, CN, CH, CS, HCN, C3, and C2H2 in the spectra
of V460 Cyg over the 2.4 to 45 m range. The theoretical models of
Helling et al. (2000) for carbon-rich AGB stars indicate
that the main contributors to molecular opacity in the 2.3
m
region, besides CO, include C2, CN, C3, and HCN. The
contribution of each to the absorption lines in our observed wavelength
band depends on the temperature and density structure, and on the C/O
ratio. Differences in these properties among our sample of carbon
stars could easily explain the variations in the observed spectra. We
have not attempted to identify other molecular lines in these spectra,
and in fact most of the individual absorption features can be
identified with CO lines. The differences in the detailed shapes of
the R-branch lines in the 2-0 band as compared with the very regular
structure in the M-giants or supergiants, however, indicates the
presence of lines from other carbon-bearing molecules. Accurate models
for the carbon star spectra in this wavelength region must include
these species in addition to CO.
We have presented spectra covering the wavelength range 2.28 to 2.36 m for a sample of 24 cool evolved stars. The sample comprises 8 M
supergiants, 5 M giants, 3 S stars, 6 carbon stars, and 2 RV Tauri
variables. The wavelengths covered include the main parts of the
12C16O v = 2-0 and 3-1 overtone bands, as well as the v = 4-2
and 13CO v = 2-0 bandhead regions. The observed transitions arise
from a wide range of energy levels extending from the ground state to
E/k > 20000 K. The spectra have a resolution of
m, or R = 3500. CO lines dominate the spectrum for all the
stars observed, and at this resolution most of the observed features
can be identified with individual CO R- or P-branch lines or blends.
We looked for correlations between the intensities of various CO
absorption line features and with other stellar properties, including
IR colors and mass loss rates. Two of the most useful CO line features
appear to be the 2-0 R14 line which is well-resolved from any other CO
or atomic features; and the CO 2-0 bandhead, which, as the shortest
wavelength of the first overtone bandheads, is not blended with any
other CO lines nor any known atomic lines. The most significant
conclusions are:
(1) The 2-0 R14 intensity shows a non-linear correlation with the 2-0 bandhead intensity, with an apparent leveling off in R14 for the deepest 2-0 bandheads (see Fig. 27).
(2) The ratio of the R14 line depth (below the continuum) to the
2-0 bandhead depth converges to a value of 0.5 for the deepest
2-0 bandheads (
I(2-0BH) <40%). Stars with shallower 2-0 bandheads
(
I(2-0BH)>50%) show a large scatter in the ratio of line depths
(Fig. 28).
(3) The published gas mass loss rates for both the AGB stars and the red supergiants in our sample correlate with the K-[12] color, consistent with other studies. The intensity of the 2-0 bandhead shows a trend with K-[12] color such that the reddest stars (K-[12] > 3 mag) exhibit a wide range in 2-0 bandhead depth, while the least reddened have the deepest 2-0 bandheads, with a small range of variation from star to star.
(4) Similarly, the reddest stars (K-[12] > 3 mag) exhibit a
large range in the ratio of the depth of the 2-0 R14 line to the 2-0
bandhead depth. Stars with less reddening (K-[12] < 3 mag) show
about one-third the range in this ratio compared to the redder sample.
If we interpret the K-[12] color in terms of a mass loss rate, the data
imply that stars with
y-1
exhibit a much narrower range in the relative strengths of CO 2-0 band
features than stars with higher mass loss rates.
(5) The wide range in CO feature intensities and intensity ratios for the most reddened stars suggests that dust continuum is not filling in the absorption, at least not in every case, but this effect needs careful modelling in any synthetic spectrum calculation.
(6) There is no evidence for a correlation of the CO band strengths with phase of the light cycle for the Miras, but intrinsic differences from star to star, or cycle-to-cycle variations, could mask any real effect. Our spectrum of AC Her, a post-AGB star, differs significantly from comparable spectra taken at other epochs. Well-sampled monitoring of selected stars is needed to determine the extent of spectral variability.
(7) Although the 12C16O overtone bands dominate the observed spectra, we note that other atomic and molecular species are present and need to be considered. First, the 13C16O 2-0 bandhead is closely blended with an atomic line of Ti I, and with the 3-1 R17 and 2-0 R0 lines of 12C16O. The observed 13C16O 2-0 bandhead depth will be affected by these features. Second, the carbon star spectra show significant modulation by other molecular lines in this wavelength region. The likely contributors include C2, CN, CH, CS, HCN, C3, and C2H2. We have not attempted to identify features of these molecules, and in fact almost all of the discrete features we observe in the carbon star spectra can be identified with the CO overtone bands. Precise modelling of this spectral region for carbon stars, however, will require that these other molecular species be included.
The range in spectral properties that we observe for this sample of cool giant and supergiant stars implies that there are significant differences in atmospheric structure. Spectra of the CO molecule should be good diagnostics for the structure of the extended atmospheres of these pulsating, mass-losing stars. Such data may be used to test hydrodynamic models which are being developed with more realistic treatment of stellar pulsations, shocks, dust formation, and molecular chemistry.
Acknowledgements
JHB thanks C. Engelbracht for invaluable help with data processing; and acknowledges support from the U.S. National Science Foundation through grants AST-9618523 and AST-9987408. We also thank the referee, Dr. M. Scholz, and Dr. J. M. Winters for helpful comments. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France.