A&A 413, 285-291 (2004)
DOI: 10.1051/0004-6361:20034136
A. F. Gulliver1, -
S. J. Adelman2,
- T. P.
Friesen1
1 - Department of Physics and Astronomy, Brandon University,
Brandon, MB, R7A 6A9, Canada
2 -
Department of Physics, The Citadel, 171 Moultrie Street,
Charleston, SC 29409, USA
Received 30 July 2003 / Accepted 5 September 2003
Abstract
We present a spectroscopic atlas of the sharp-lined, hot
metallic-line star o Pegasi (A1 IV) based on spectrograms
obtained with the long camera of the 1.22-m telescope of the
Dominion Astrophysical Observatory using a Reticon detector. For
3826-4882 the inverse dispersion is 2.4
Å mm-1 with a resolution of 0.072 Å. At
the continuum the mean signal-to-noise ratio is 800. The
wavelengths in the laboratory frame, the equivalent widths, and
the identifications of the various spectral features are given.
For studies of similar stars and for atomic physicists interested
in improving atomic line parameters, this atlas should provide
useful guidance. The stellar and synthetic spectra with their
corresponding line identifications can be examined at
http://cdsweb.u-strasbg.fr/cgi-bin/qcat?/A+A/413/285.
Key words: atlases - stars: early-type - stars: individual - o Pegasi - stars: chemically peculiar
o Pegasi (43 Pegasi, HR 8641, HD 214994, BD +284436, HIP 112051) is a sharp-lined and unreddened prototype of the
hot metallic-line stars. Recently Adelman et al. (2003) found that its
evolutionary track indicates that during its lifetime in the main
sequence band when it was closer to the Zero Age Main Sequence, it
was a Mercury-Manganese star. Its spectral type of A1 IV by
Cowley et al. (1969) confirms both Osawa (1959) and
Ljunggren & Oja (1961).
The spectroscopic material used for this atlas was analyzed by Adelman et al. (2003) who provide information on previous studies of the optical region for this star. It is assembled from a set of very high quality spectra obtained with a Reticon on a spectrograph with a known amount of scattered light which has been removed. By making widely available this spectrum in FITS and HTML formats and its measurements, we hope this material will be useful for other studies.
This spectral atlas of o Pegasi's spectrum
3826-4882 is based on 2nd order exposures obtained
with a Reticon detector and the IS96B image slicer with wavelength
coverage of 67 Å at the long camera of the 1.22-m telescope of
the Dominion Astrophysical Observatory (DAO) and with a resolution
of 0.072 Å (two pixels) or a resolving power of 60 000. Its mean
signal-to-noise (S/N) ratio is 800. The central wavelengths of
the 19 spectrum sections between
3860 and
4850
were separated by 55 Å allowing a several Å overlap between
adjacent sections. A central stop placed in the beam removed light
in the same manner as the secondary mirror of the telescope. The
exposures were flat fielded with those of an incandescent lamp
placed in the coudé mirror train as viewed through a filter.
Reticon exposures were reduced to one-dimensional FITS files with the Program RET72 (Hill & Fisher 1986) utilizing the lamp exposures. Flat images were summed to create mean flat exposures. The arc and stellar exposures were then divided by the mean flat image.
We measured the arc files interactively using the relevant routine in the spectrophotometric reduction and analysis program REDUCE (Hill & Fisher 1986). An initial approximation to the dispersion characteristics of each DAO spectrograph were input. Then corrections, based on the agreement between the predicted and measured position of each line, were used to predict the position of each new line. When the corrections became sufficiently small, the remaining lines were measured automatically. Heliocentric radial velocity corrections were calculated with the program VSUN (Hill & Fisher 1986). Wavelength-calibrated spectra were produced with REDUCE using the arc files. The wavelength scale accuracy is better than 0.005 Å.
Table 1 lists the spectrograms with their exposure numbers, Heliocentric Julian Dates at the mid-points of their exposures, central wavelengths, the derived radial velocities and their associated errors, and an estimate of the S/N ratio. The S/N ratio of each section was estimated from the root-mean-square deviation for the continuum point intervals, usually smooth regions without lines close to the continuum, as part of the rectification process. The mean of these S/N ratios is about 800:1.
Table 1: o Pegasi spectrograms.
The radial velocity of each spectrum was measured using the
program VCROSS (Hill & Fisher 1986) that cross-correlated the
stellar with a synthetic spectrum calculated with the preliminary
atmospheric parameters of
K, log g= 3.6, 0.2 dex times solar metallicity, a microturbulence of 1.7 km s-1and a V
of 7 km s-1, produced by the program SYNTHE
(Kurucz & Avrett 1981). The cross-correlation function was fitted
by a Gaussian, the centroid and FWHM of which were allowed to
vary. A zero background slope of the Gaussian fit was an important
restriction. The mean error of the radial velocities, as shown in
Table 1, is 0.3 km s-1. All spectra were shifted to rest
wavelengths before further processing.
The stellar intensity files were rectified with REDUCE so that the
continuum was calculated from locally averaged points. The
rectification was completed by interpolating between the averaged
data with Hermite spline functions, which always pass through the
averaged continua at the selected wavelengths. The scattered light
along the spectrum was assumed to be 3.5% of the continuum
(Gulliver et al. 1996). The atlas and the published equivalent widths
reflect this correction. This initial rectification was performed
by the usual choosing of suitable continuum points by visual
inspection. The resultant rectified spectrum served as the basis
for the analysis of the spectrum reported in Adelman et al. (2003).
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Figure 1:
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Open with DEXTER |
The final stellar parameters of o Pegasi were determined,
as described in Adelman et al. (2003), by fitting the initial
rectification version of the observed spectrum including an
extracted H
profile plus the relative continuous fluxes
from both the IUE UV and visible as tabulated by Adelman et al. (1989).
The parameters found were
K, log
,
the individualized abundances of
Adelman et al. (2003), a microturbulence of
km s-1and a V
of
km s-1. The parameter
determination used synthetic spectra of the H
regions from
ATLAS9 LTE plane parallel model atmospheres (Kurucz 1993) with
Program SYNTHE (Kurucz & Avrett 1981) as well as the predicted
fluxes with ATLAS9 for comparison with the observations. The
and log g were determined by simultaneously fitting
the continuous flux including both IUE UV and visible and the
extracted H
profile using STELLAR (Hill et al. 1996;
Hill et al. 2003). The H
profile was predominantly
sensitive to log g and the shape of the continuous fluxes to
.
V
was determined by fitting the spectrum from
4450-4800 using STELLAR. The microturbulence was
determined from the fine analysis of Adelman et al. (2003).
These stellar parameters were used in turn in the program, STELLAR, to generate the synthetic spectrum that was convolved with a digitally sampled instrumental profile with a FWHM of 0.072 Å. The final rectification points for the atlas of o Pegasi were then chosen by a novel technique that involves the selection of suitable intervals from the synthetic spectrum of o Pegasi. From this synthetic spectrum rectification points can be selected that need not be actual continuum points. A trial choice of suitable rectification points is based upon any wavelength interval for which the synthetic spectrum is roughly smooth and there is good agreement with the observed spectrum. Obviously, good stellar parameters are a necessity. In effect this is an iterative process in which trial points are modified or rejected if there is not good agreement between the synthetic and observed spectrum. The process is complicated by poorly known atomic data for some lines which produce discrepant line strengths and positions in the synthetic spectrum. For any suitable rectification point, an intensity level is established from the synthetic spectrum and the observed spectrum is normalized at that value.
This technique can be used for any star that has well defined stellar parameters. It is particularly useful for rectification across broad hydrogen wings, for late type stars in which true continuum may be entirely absent and, as in this case, the appending of sections of spectra to produce a single monolithic spectrum. To facilitate the seamless combination of the sections there were at least two common rectification points over each several Å overlap.
The final monolithic spectrum is displayed at the URL http://www.brandonu.ca/physics/gulliver/atlases.html. The JAVA tool provided there allows the display and comparison, in any combination, of the observed and synthetic spectra and their respective line identifications. Unrectified and rectified sections of the o Pegasi spectrum, and the complete rectified spectrum, are also available as FITS format files from the first author (AFG) and the CDS. Copies of the line identifications are also available.
For the purpose of illustrating the nature of the atlas, Figs. 1
and 2 show two adjacent 10 Å pieces,
4450-4460
and
4460-4470 of the section centered at 4465 Å
of the o Pegasi atlas, which includes 19 such
sections. Although the majority of the line identifications for
this section are reproduced in the figures, multiple possible
identifications of a given feature are not included to avoid
overcrowding. The Java tool mentioned above gives a more accurate
impression of the quality of the o Pegasi atlas.
Figures 1 and 2 also include the final synthetic spectrum produced by STELLAR as described above. The differences between the observed and synthetic spectra are striking, clearly illustrating the deficiencies in the atomic line parameters. This atlas and others can be used to improve these values.
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Figure 2:
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Open with DEXTER |
We employed the program VLINE (Hill & Fisher 1986) to measure for
each line the equivalent width, the central wavelength, the line
depth, and the full width at half maximum of the fitted profile,
which was taken to be a Gaussian for metal lines except for some
He I lines which were Lorentzian profiles. Our rotational
velocity estimate based on non-blended lines near Mg II
4481 was 6 km s-1. In measuring the spectrum, we
used the fixed width profile feature for weak lines and to
deconvolve blended lines requiring that their widths correspond to
our derived rotational velocity estimate.
To begin the line identification process, we identified the cleanest lines in the spectrum which are minimally affected by noise and by blending components. These can often be found by examining stellar line identification lists of stars of similar temperature, previous studies of o Pegasi, or working with standard references. As not all atomic wavelength studies have equally well-determined wavelengths, we preferred to use those whose values are consistent with modern interferometric determinations for Fe I and Fe II. Stellar lines were first identified with the general references A Multiplet Table of Astrophysical Interest by Moore (1945), Wavelengths and Transition Probabilities for Atoms and Atomic Ions, Part 1 by Reader & Corliss (1980), and selected references from the bibliography of Adelman & Snijders (1974) whose most recent update is Adelman (2001). We used line identifications by Adelman and his associates for other stars which they have analyzed using DAO spectrograms in previous papers of this series.
Table 2:
Line measurements of o Pegasi
(
4450-4460 Region).
A sample of the line identifications is presented in Table 2 for
the
4450-4470 section. The full identifications
are available elsewhere as noted previously. To identify as well
as possible the lines in the spectrum of o Pegasi after the
elemental abundances of o Pegasi were initially determined,
a synthetic spectrum was calculated using Program SYNTHE
(Kurucz & Avrett 1981) with the adopted model atmosphere, solar
abundances, the atomic data of Kurucz & Bell (1995), the
instrumental profile of the long camera of the DAO coudé
spectrograph, and the other parameters initially found for
o Pegasi by Adelman et al. (2003). It was a good, but not
perfect match. A list of lines which contributed significantly to
the spectrum was made and used to help identify particularly the
unidentified features. The major changes were the identification
of Co II lines and some additional Fe II and Ni II lines as
described in Adelman et al. (2003). The synthesized spectrum which
is provided with this atlas was calculated after these additional
clean lines which were found to be present were used to improve
the derived abundances while those lines which were initially used
and found to be blended were removed.
For those lines not in Moore (1945), letters are used in place of multiplet numbers to indicate the other sources: C = Catalan et al. (1964), D = Dworetsky (1971), G = Guthrie (1985), H = Hudlt et al. (1982), I = Iglesias et al. (1988) for V II and Iglesias & Velasco (1964) for Mn II, J = Johansson (1978), K = Kiess (1951),Kiess (1953), KX = Kurucz & Bell (1995), L = Litzen (2002), MCS = Meggers et al. (1975), N = Nilsson et al. (1991), and P = Pettersson (1983). Some multiplet numbers are from Moore (1965) for Si II and from Moore (1993) for C I and O I.
In Table 2 the far left column contains the letters B and R,
standing for blue and red, which are guides to the range of lines
whose wings are at least somewhat blended together. For example
in Fig. 1, the features corresponding to N I 4214.804 and Fe I 4216.1838 are labelled B and R, respectively, because they are
blended with the shortward and longward wings of the Sr II 4215.524, Fe I 4215.4777 blend. The remaining columns are the
laboratory wavelength in Å (the stellar wavelength as corrected
for the stellar radial velocity of each spectrum), the equivalent
width (
)
in mÅ, the line depth as a fraction of the
continuum height, the line width (FWHM) in Å, and the identified
atomic lines which cause the observed feature. The stellar and the
laboratory wavelengths should be close, but blending and errors
can produce discrepancies. Possible identifications are given in
parentheses and brackets indicate that an identified line may be
contributing to two measured stellar features.
The following discussion indicates which atomic species were found in the observed spectral range of o Pegasi. A Multiplet Table of Astrophysical Interest by Moore (1945) was the primary source of line identifications. When other references were used which substantially replaced or supplemented this source they are given immediately after the species name. In general species not identified are not discussed.
1. H I - The Balmer lines are present.
2. He I - The strongest and medium strength He I lines in the region studied are present.
3. C I - Moore (1993) - The stronger lines of multiplet 6 are cleanly present while the weaker members are blended. One line of multiplet 5 is also cleanly present.
4. C II - Moore (1993) - One line of multiplet 6 is unblended while the other is blended.
5. O I - Moore (1993) - Lines of multiplets 3 and 5 are present.
6. Mg I - Lines of multiplets 3, 11, 14, 15, and 17 are cleanly present while that of multiplet 16 is blended.
7. Mg II - Lines of multiplets 4, 5, 9, 10, and 18 are present.
8. Al I - The two lines of multiplet 1 are present.
9. Al II - 4663.054, multiplet 2, is present and
3900.680,
multiplet 1, is probably part of a blend.
10. Si I - Moore (1967) - 3905.523, multiplet 3, the
strongest line in the observed region is present.
11. Si II - Moore (1965) - Lines of multiplets 1, 3, 3.01, 7.05, 7.06, 7.26, and 20 are present.
12. S I - The three lines of multiplet 2 are present.
13. S II - Pettersson (1983) - All lines with intensities 20 are present as is one intensity 19 line and probably a few
other weaker lines.
14. Ca I - Lines of multiplets 2, 4, 5, 6, 23, 25, 37, and 51 are present.
15. Ca II - The H and K lines of multiplet 1 are present.
16. Sc II - Lines of intensity 8 of multiplets 7, 8, 14, and 24 are
present as well as a few weaker lines with intensities of 2 or more.
17. Ti I - Almost all lines with intensities 60 are present as are
many with intensities between 35 and 55.
18. Ti II - Litzen (2002) - All lines with intensities 64
are present as are many with intensities between 6 and 64 and some
lines from Hudlt et al. (1982) and from Moore (1945).
19. V I - The ultimate line 4379.24 probably corresponds to a 1.8 mÅ
feature.
20. V II - Iglesias et al. (1988) - Almost all lines with intensities
20 are present as are a few weaker lines. A few lines only
from Moore (1945) are retained.
21. Cr I - Kiess (1953) - Almost all lines with intensities
100 are present along with many with intensities between 75
and 100. A few weaker lines are marked as possible
identifications.
22. Cr II - Kiess (1951) - Almost all lines of intensity 4
are present as well as some with intensities between 1 and 3. Some
lines from Dworetsky (1971) supplement this source.
23. Mn I - Catalan et al. (1964) - All lines with intensities 2000 are present as well as most with intensities
500 and a
few with intensities
200.
24. Mn II - Iglesias & Velasco (1964) - Almost all lines with intensities greater than 40 are present as well as most with intensities between 20 and 40.
25. Fe I - Most lines with intensities 1 are found along with many with
intensities in parentheses and some formerly predicted lines from Nave et al.
(1994). The wavelengths when available are from Nave et al. (1994).
26. Fe II - Johansson (1978) - Many lines from
Dworetsky (1971) and Guthrie (1985) are present. All lines
from Johansson's Table I with intensities 3 and some
intensity 1 and 2 lines are present except close to the core of
H
.
All lines from his Table II are present. As noted by
Adelman (1987) many predicted Fe II lines in Moore (1945) are
present.
27. Fe III - Glad (1956) - Only 4419.599, the strongest
line and
4431.015, one of the next strongest lines, of
multiplet 4 are found.
28. Co I - Lines with intensities 30 are present along with many of
intensities 20 and 25.
29. Co II - Kurucz & Bell (1995) - The synthetic spectrum of o Pegasi with the initial Co abundance from the Co I lines shows that several Co II lines are present.
30. Ni I - Most lines of intensity 7 are present. Many with intensities
of 2 to 6 are present.
31. Ni II - Lines of multiplets 9, 10, 11, 12, and 13 are present as well as some not included in Moore (1945) according to the synthetic spectral calculations.
32. Zn I - The three lines of multiplet 1 are present.
33. Cd I - There are features close to the position of two lines of multiplet
2 4799.918 and
4678.160. In other stars of similar
temperature both yield very large overabundances and so are marked only as
possible identifications. The former might be due to instead to a Ca II line.
34. Sr II - The two strong lines of multiplet 1 are present as well the two lines of multiplet 3, one of which is blended.
35. Y II - Nilsson et al. (1991) - All lines with intensities 165
are present as well as many with intensities between 43 and 164.
36. Zr II - Almost all lines with intensities 3 are present
as well as many intensity 1 and 2 lines.
37. Ba II - Lines of multiplets 1, 3, and 4 are present, some of which are blended.
38. La II - Meggers et al. (1975) - Many lines are present with
intensities 1100.
39. Ce II - Meggers et al. (1975) - Many lines are present with
intensities 860 along with a few weaker ones.
40. Nd II - Meggers et al. (1975) - Of the five strongest lines, 3 are clearly present and 2 are parts of blends.
41. Eu II - Meggers et al. (1975) - Several of the strongest lines are present.
42. Gd II - Meggers et al. (1975) - Of the five strongest lines, two are clearly present, two are blended, and one is absent. Gd II is regarded as being weakly present.
43. Tb II - Meggers et al. (1975) - Only the strongest line in the
observed region, 3948.73, is a probable identification of
a 3.1 mÅ feature.
44. Dy II - Meggers et al. (1975) - Of the strongest three lines in the region present, one is cleanly identified and the other two are blended. This species is regarded as having lines which are weakly present.
45. Ho II - Meggers et al. (1975) - The strongest line in the observed
region, 4045.44, might be blended with Co I(41)
4045.386, but was not marked as a definite
identification in the line list.
46. Er II - Meggers et al. (1975) - The strongest line in the observed
region, 3906.31, might contribute to a 4.8 mÅfeature.
47. Tm II - Meggers et al. (1975) - The strongest line in the observed
region, 3848.02, might contribute to a 1.8 mÅ feature
along with a weak Y II line.
48. Yb II - Meggers et al. (1975) - The strongest line in the region
observed, 4180.81, might be the identification of a 2.0 mÅ feature.
Acknowledgements
AFG and SJA thank Dr. James E. Hesser, Director of the Dominion Astrophysical Observatory for the observing time. This research was supported in part by grants from the National Sciences and Engineering Research Council of Canada and The Citadel Foundation.