A&A 394, 187-192 (2002)
DOI: 10.1051/0004-6361:20021117
H. Caliskan1 - S. J. Adelman2,3 - M. T. Cay1 - I. H. Cay1 - A. F. Gulliver3,4 - G. H. Tektunali1 - D. Kocer1 - A. Teker1
1 - Department of Astronomy and Space Sciences, Istanbul University,
34452 University, Istanbul, Turkey
2 - Department of Physics, The Citadel, 171 Moultrie Street, Charleston, SC 29409, USA
3 - Guest Investigator, Dominion Astrophysical Observatory, Herzberg
Institute of Astrophysics, National Research Council of Canada,
5071 W. Saanich Road, Victoria V9E 2E7, Canada
4 - Department of Physics, Brandon University, Brandon, MB R7A 6A9, Canada
Received 4 September 2001 / Accepted 1 August 2002
Abstract
Elemental abundances analyses of the superfically normal stars
Ser
(F6 V) and 101 Herculis (A7 V) consistent with previous studies of this series
using photographic region spectrograms obtained with Reticon and CCD detectors
produced derived values that are generally solar except for Al which is very
underabundant and the rare earths which do not have the solar pattern. Similar
discrepancies from solar values are seen in other superficially normal stars of
this series with similar temperatures. Our results for
Ser are in
acceptable agreement with other recent studies based on different techniques and
spectral regions.
Key words: stars: abundances - stars: individual:
Ser - stars:
individual: 101 Her
Paper XXIII (Adelman et al. 2000), which investigated 28 And (A7 III) and 99 Her (F7 V), was our most recent paper about superficially
normal late A and F type stars. By investigating two somewhat similar stars,
Ser (F6 V) and 101 Herculis (A7 V), we increase the number of such
objects consistently analyzed. In the photographic region placing their
continuua can be especially difficult as in the blueward
part this depends on isolated windows. This task is greatly aided by high
resolution high signal-to-noise spectra. Figure 1 shows an illustrative portion
of these spectra. Many of the lines in
Ser can be seen to be blended
in 101 Her.
![]() |
Figure 1:
The normalized spectra of 101 Her and ![]() |
Open with DEXTER |
The abundances of
Ser (HD 142860, HR 5933), spectra type F6 V (Gray et al. 2001a), have been determined in many studies of invididual elements, e.g.,
Co by Adelman et al. (2000), and in several comprehensive studies such as
those by Chen et al. 2000 who found
K,
,
and
,
Gray et al. (2001b) who found
= 6350 K,
,
,
and Carretta et al. (2000) who found
K,
,
,
and
.
Earlier Fuhrmann
(1998) included
Ser among the some fifty nearby F- and G-stars. He
considered it a suspected binary, and gave
K,
,
,
and v sin i = 10.6 km s-1. This well
studied F star permits us to compare our results with those of solar type
stars often analyzed relative to the Sun.
101 Her (HD 166230, HR 6794), spectral type A7 V and
km s-1
(Abt & Morrell 1995), is not particularly well studied. Jerzykiewicz (1993)
found it was probably photometrically constant while Adelman (2001) found it
was one of the least variable stars according to Hipparcos photometry. Earlier
classifications indicated it as a giant (see, e.g. Cowley et al. 1969).
Star |
![]() |
Log g | Method |
![]() |
6302 | 3.56 | Napiwotzki et al. (1993) with uvby![]() |
6300 | 4.25 | Spectrophotometry and H![]() |
|
6300 | 5.00 | Spectrophotometry and H![]() |
|
6300 | 4.00 | H![]() |
|
101 Her | 8091 | 3.44 | Napiwotzki et al. (1993) with uvby![]() |
8061 | 3.51 | Photometric values corrected for offset found by Adelman et al. (2002), Mixing Length theory | |
8061 | 3.69 | Photometric values corrected for offset found by Adelman et al. (2002), CM theory |
Number | ![]() |
![]() |
|||||
Star | Species | of Lines | (km s-1) |
![]() |
(km s-1) |
![]() |
gf values |
![]() |
Fe I | 297 | 1.2 |
![]() |
1.4 |
![]() |
MF+KX |
250 | 1.2 |
![]() |
1.4 |
![]() |
MF | ||
Fe II | 36 | 1.1 |
![]() |
1.1 |
![]() |
MF+KX | |
adopted | 1.2 | ||||||
101 Her | Fe I | 120 | 4.4 |
![]() |
4.4 |
![]() |
MF+KX |
107 | 4.4 |
![]() |
4.4 |
![]() |
MF | ||
Fe II | 26 | 4.6 |
![]() |
4.6 |
![]() |
MF+KX | |
adopted | 4.5 |
gf value references: MF = Fuhr et al. (1988), KX = Kurucz & Bell (1995).
Note: For ![]() ![]() equivalent widths and have minimum scatter about the mean, respectively. |
We fit Gaussian profiles through the spectral lines of
Ser. For 101 Her
lines with equivalent widths
100 mÅ required rotational
profiles, with equivalent widths
14 mÅ Gaussian profiles, and
with intermediate equivalent widths the profile which fit best. Rotational
velocity estimates from clearly single medium strength lines near
4481 are 9 km s-1 for
Ser and 41 km s-1 for 101 Her. Fekel (1997) finds
km s-1 and
km s-1 for
Ser. The stellar lines were identified with
the general references A Multiplet Table of Astrophysical Interest (Moore 1945)
and Wavelengths and Transition Probabilities for Atoms and Atomic Ions, Part
1 (Reader & Corliss 1980) as well as Huldt et al. (1982) for Ti II, Catalan
et al. (1964) for Mn I, Iglesias & Velasco (1964) for Mn II, Nave et al.
(1994) for Fe I, and Johansson (1978) for Fe II.
For
Ser, clean lines of C I, Na I, Mg I, Mg II, Al I, Si II, Ca I,
Sc II, Ti I, Ti II, V I, V II, Cr I, Cr II, Mn I, Mn II, Fe I, Fe II, Co I,
Ni I, Ni II, Zn I, Sr I, Sr II, Y II, Zr II, Ba II, La II, Ce II, Pr II, Nd II,
Sm II, Eu II, and Gd II were found. In addition lines of Si I were blended
while it was difficult to measure the very strong Ca II K and H lines. There
are also molecular lines which were not studied.
In the spectrum of 101 Her, we used lines of O I, Mg I, Mg II, Al I, Si II, Ca I, Sc II, Ti II, V II, Cr I, Cr II, Mn I, Mn II, Fe I, Fe II, Co I, Ni I, Ni II, Zn I, Sr II, Y II, Zr II, Ba II, La II, Nd II, Sm II, Eu II, and Gd II in our elemental abundance analysis. There are also Ca II and lines of a few other species which are probably in blends such as Si I and Ti I.
We derived the radial velocities from comparisons of the stellar and laboratory
wavelengths after corrections were applied for the Earth's orbital velocity.
For 18 spectrograms, our mean radial velocity was
km s-1
for
Ser while Abt & Biggs (1972) list values between 2.6 and 7.8 km s-1. As the standard deviation of the mean for the average value is
not much larger than the standard deviation of the means for individual
spectra of order 0.60 km s-1, at best this is weak evidence for
radial velocity variability. For 101 Her, we found a mean value of
km s-1 for 16 site 2 spectrograms which compares with values
between -12.2 and -19.8 km s-1 compiled by Abt & Biggs (1972). Thus 101
Her is a candidate to be studied further for possible binarity. It might be
possible to reduce the errors in the radial velocity by cross correlating
synthetic with observed spectra.
Table 1 lists our effective temperature and surface gravity estimates with the
last values for each star being those adopted. To get initial estimates from
the homogeneous uvby
data of Hauck & Mermilliod (1980, 1998), we
used the computer program of Napiwotzki et al. (1993). The uncertainties are
about
200 K and
0.2 dex (Lemke 1989). To refine these values for
Ser we calculated synthetic spectra of the H
region 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 by Schild (unpublished prior to
being included in Breger 1973). As Smalley & Kupka (1997) argued that the
turbulent convection theory of Canuto & Mazzitelli (1991, 1992) should be more
realistic than Mixing Length theory (Castelli et al. 1997) and as
Kupka (private communication) supplied the necessary subroutine for
implementing Canuto-Mazzitelli convection in ATLAS9, we have used these models
in the determination of the effective temperature and surface gravity. Smalley & Kupka (1997) found
K,
.
The
corresponding values for Mixing Length theory are also given. Our effective
temperature is in the middle of the range of recent determinations, but our
surface gravity is much larger the average. But we found that Fe I and Fe II
lines gave substantially different abundances for this surface gravity. As
our photometrically determined value of
was close to the mean of
many recent determinations, we adopted it. Then to find
we demanded
ionization equilibrium from Fe I and Fe II lines and considered
the results for other elements with lesser weight, a process which indicated
.
Our final values are close to those found using the infrared
flux method by Smalley & Kupka (1997)
K,
and also by Alonso et al. (1996)
K.
Species | 99 Her | ![]() |
28 And | 101 Her | Sun | ||
log N/H | log N/H | Lines | log N/H | log N/H | Lines | ||
C I | -3.51 |
![]() |
2 | -3.59 | ... | 0 | -3.45 |
O I | ... | ... | 0 | ... | -3.65 | 1 | -3.13 |
Na I | -5.95 |
![]() |
2 | ... | ... | 0 | -5.67 |
Mg I | -5.01 |
![]() |
2 | -4.60 |
![]() |
2 | -4.42 |
Mg II | -4.25 |
![]() |
2 | -4.52 |
![]() |
4 | -4.42 |
Al I | -6.65 | -6.01 | 1 | -6.35 | -6.13 | 1 | -5.53 |
Si II | -4.81 |
![]() |
2 | -4.71 |
![]() |
2 | -4.45 |
Ca I | -6.07 |
![]() |
12 | -5.89 |
![]() |
12 | -5.64 |
Sc II | -8.92 |
![]() |
6 | -8.99 |
![]() |
2 | -8.83 |
Ti I | -7.39 |
![]() |
41 | -7.12 | ... | 0 | -6.98 |
Ti II | -7.22 |
![]() |
26 | -7.10 |
![]() |
25 | -6.98 |
V I | -8.24 |
![]() |
17 | -8.15 | ... | 21 | -8.00 |
V II | -8.01 |
![]() |
9 | -8.05 |
![]() |
5 | -8.00 |
Cr I | -6.61 |
![]() |
85 | -6.52 |
![]() |
7 | -6.33 |
Cr II | -6.51 |
![]() |
18 | -6.34 |
![]() |
12 | -6.33 |
Mn I | -7.04 |
![]() |
27 | -6.85 |
![]() |
8 | -6.61 |
Mn II | -6.61 | ... | 0 | -6.24 | -6.38 | 1 | -6.61 |
Fe I | -4.94 |
![]() |
297 | -4.71 |
![]() |
120 | -4.50 |
Fe II | -4.76 |
![]() |
36 | -4.64 |
![]() |
26 | -4.50 |
Co I | -7.31 |
![]() |
12 | -7.57 |
![]() |
2 | -7.08 |
Ni I | -6.22 |
![]() |
43 | -6.04 |
![]() |
8 | -5.75 |
Ni II | -5.78 |
![]() |
2 | -5.73 |
![]() |
2 | -5.75 |
Zn I | -7.84 |
![]() |
2 | -7.84 |
![]() |
2 | -7.40 |
Sr I | -8.94 | -9.03 | 1 | ... | ... | 0 | -9.03 |
Sr II | -9.61 |
![]() |
3 | -8.95 | -9.20 | 1 | -9.03 |
Y II | -10.02 |
![]() |
7 | -9.85 | -10.22 | 1 | -9.76 |
Zr II | -9.53 |
![]() |
8 | -9.51 |
![]() |
8 | -9.40 |
Ba II | -10.22 | -9.99 | 1 | -9.77 | -9.91 | 1 | -9.87 |
La II | -10.42 |
![]() |
9 | -10.87 |
![]() |
3 | -10.83 |
Ce II | -10.37 |
![]() |
24 | -10.20 | ... | 0 | -10.42 |
Pr II | -11.51 |
![]() |
2 | ... | ... | 0 | -11.29 |
Nd II | -10.61 |
![]() |
14 | -10.62 | -9.59 | 1 | -10.50 |
Sm II | -10.58 |
![]() |
13 | -10.53 | -9.84 | 1 | -10.99 |
Eu II | -11.87 |
![]() |
3 | -11.33 | -10.31 | 1 | -10.49 |
Gd II | -10.22 |
![]() |
8 | -10.91 | -9.98 | 1 | -10.88 |
![]() |
6100 | 6300 | 7350 | 8061 |
For 101 Her, which does not have good spectrophotometric values, we correct
the photometric values using the mean offsets found by Adelman et al.
(2002). The surface gravity is perhaps slightly smaller than one would expect
for a main sequence star. We used SYNTHE (Kurucz & Avrett 1981) to compute
the H
region and found that the observed and theoretical H
profiles agree well. This confirms that the effective temperature
is well determined. For the best determined microtubulence following the
method described below using 2 km s-1 odfs (opacity distribution
functions), the difference between the iron abundances derived from Fe I and
Fe II lines is 0.09 dex which is acceptable. But as the microturbulence was
4.4 km s-1 we calculated a new model with the same parameters and the
odfs for 4 km s-1 which further improves the agreement.
The metal abundances were determined using program WIDTH9 (Kurucz 1993) with
line damping constants from Kurucz & Bell (1995) or semi-classical
approximations in their absence. Abundances from Fe I and Fe II lines were
derived for a range of possible microturbulences whose adopted values (Table 2)
result in the derived abundances being independent of the equivalent widths
()
or having a minimal scatter about the mean (
)
(Blackwell et al. 1982).
For
Ser, the derived mean abundances from Fe I and Fe II lines agree
when
= 1.2 km s-1. This value is close to 1.5 km s-1 of Gray
et al. (2001b). For 101 Her, we find
= 4.5 km s-1 with the Fe I and
Fe II line results in excellent agreement. We assumed that the helium
abundances were solar. Thus to convert
values to log N/H values
-0.04 dex were added.
Table 3, the analyses of the line spectra, contains for each line the
multiplet number (Moore 1945), the laboratory wavelength, the logarithm of the
gf-value and its source, the equivalent width in mÅ as observed, and the
deduced abundance. Source references are given at the end of this table.
Letters are used in place of multiplet numbers to indicate sources other than
Moore (1945): C = Catalan et al. (1964), D = Dworetsky (1971), I = Iglesias &
Velasco (1964), J = Johansson (1978),
& Bell (1995), and N = Nave et al. (1994).
This study's abundances are compared with those of the Sun (Grevesse et al. 1996) in Table 4. Also given are the number of lines and the abundances of 28 And and 99 Her (Paper XXIII) whose values in general suggest that they are least slightly metal-poor compared to the Sun if 0.3 dex is considered to be a measure of significance. All Al abundances are significantly less than solar . It would be useful to study other lines of Al I to see if they also show this effect. Although 101 Her has greater rare earth abundances than the other three stars, those based on only a few lines need to be regarded with suspection. For the other three stars that the derived Eu values are so underabundant is somewhat surprising.
The 101 Her results for Nd, Sm, and Gd, which are based on only one line each,
that suggest some rare earth elements are greatly overabundant need
additional confirmation. Of the remaining 19 non-rare earth elements the mean
underabundance of 101 Her with respect to the Sun is
dex which
is solar. The major exception is Al which is underabundant by -0.60 dex. We
used the Mg II value rather that than for Mg I as the lines of the later
species are in very heavily line blanketed regions. The values for O, Sr, Y,
Ba, and Eu which depend on only one line each need confirmation. Cr I and
Cr II lines give somewhat similar results while those from Ni I and Ni II
lines are not in quite as good agreement. This is a difficult star to analyze
due to its moderate rotation and degree of line blanketing.
For 26 species the mean overabundance for
Ser is
dex
which is solar using the Mg I value as it is more reliable than that from Mg
II. Al is underabundant as are Si and Eu while Sm and Gd are overabundant.
The mean and the iron abundances have solar values rather than being
marginally underabundant as most other recent analyses. Table 5 compares the
results of this study with those of Caretta et al. (2000), which are a
renalysis of Edvardsson et al. (1993), Chen et al. (1998), and Nissen et al.
(2000). Our microturbulence is the smallest while our Fe abundance which
depends in part on strong lines is the largest. The other three values in
common with those from Caretta et al. are in good agreement. But with Chen et al. and Nissen et al. in some cases our agreement is not as good. The [Si/Fe]
values have a discrepancy of 0.41 dex. Although our Si II lines have good gf values and are not subject to nLTE effects as are many Si II lines in the red,
they are in a part of the spectrum strongly affected by line blanketing and so
have relatively large errors.
Quantity | Caretta | Chen or | This |
et al. | Nissen et al. | Study | |
[Fe/H] | -0.18 | -0.22 | -0.04 |
[C/Fe] | -0.05 | ... | -0.09 |
[Na/Fe] | 0.02 | -0.06 | -0.06 |
[Mg/H] | 0.07 | ... | -0.03 |
[Si/Fe] | ... | 0.07 | -0.34 |
[Ca/Fe] | ... | 0.11 | -0.12 |
[Ti/Fe] | ... | -0.12 | -0.14 |
[Sc/H] | ... | 0.25 | 0.01 |
[Mn/H] | ... | 0.01 | -0.08 |
![]() |
6268 | 6227 | 6300 |
![]() |
4.04 | 4.18 | 4.00 |
![]() |
1.63 | 2.15 | 1.20 |
An important question concerns the uncertainties in the results for
Ser. Canonically the smallest errors for absolute abundance analyses are
thought to be about 0.20 dex (see, e.g., Gigas 1988). One can perform a
sensitivity test of the abundances. Table 6 shows the changes in the 14
derived abundances based on the most lines for
Ser. We considered the
effective temperature being increased by 200 K, the surface gravity by 0.2 dex,
the equivalent width scale by 10%, and the microturbulence by 0.2 km
s-1. For each species the total error is the square root of the square
sum of the individual errors (see, e.g. Lemke 1990). These values range from
0.11 to 0.22 dex and are similar to the errors in Table 4. Differences in
sensitivity are species and mean line strength dependent. In a homogeneous
set of analyses for many similar stars, one can select those lines which give
the most consistent results and thus reduce the resultant errors a little.
The offset in [Fe/H] of this study for
Ser relative to those of
previous studies is of concern. If we use Fe II lines only from
Fuhr et al. (1988), the mean abundance is
,
a reduction
of 0.08 dex. This suggests a small systematic offset in the Fe II gf values
from Kurucz & Bell (1995) relative to those in Fuhr et al. (1988).
Using only the 10 Fe I lines whose equivalent widths are between 10 and 30
mÅ and whose gf values have errors less than 25% in Fuhr et al.
(1988),
,
a reduction of 0.06 dex. Thus about 0.07 dex of the offset may be due to gf values. The lack of substantial reductions
of the scatter in these Fe I and Fe II determations relative to those of the
entire sets of Fe I and Fe II lines also suggests that part of the scatter in
each elemental abundance value is due to systematic errors in the gf values.
As the absolute gf values of the lines in the photographic are often better determined than those further to the red, the above paragraph not with standing, this is possibly a source of problems in the comparison with other studies. The results for certain species may include possible nLTE effects. But our continuum is less certain and the degree of undetected line blending probably greater than in the red. For a proper comparison with our other stars, we need to study the photographic region as many species have lines there. At our S/N ratio of 200+ and resolution of 0.072 Å (2 pixels) we reliably interpolate the continuum in the blue even in late F stars. The line density dramatically decreases as one proceeds longward of about 4630 Å at which point determining the continuum level also becomes easier. Thus studies of F type stars beginning about 4630 Å and going longward should have most of the virtues claimed for studies in the red, especially better determined continua and a sufficiently low line density so that line blending is not a major problem. Further the reduction of echelle spectrograms involve special problems including the ripple which do not affect coudé spectrograms. By analyzing the same lines as other investigators as well as those of this study, we should be better able to identify and then hopefully to eliminate sources of possible difference, including errors in the gf values. It is probably also desirable to study another well studied late F star using spectra which include the red region. However, some of the systematic errors discussed by Kurucz (2002) for solar type star studies may dominate any comparison.
In summary, our abundance study of
Ser based on photographic region
spectrograms shows decent agreement with those of other investigators who used
spectra further to the red. To understand the origins of the discrepancies, it
is desirable to obtain spectra of this star in the red and then analyze them
consistently with those used in this paper.
Species |
![]() |
![]() |
![]() |
![]() |
Total |
Ca I | 0.14 | 0.00 | 0.17 | 0.02 | 0.22 |
Ti I | 0.15 | 0.01 | 0.07 | 0.03 | 0.17 |
Ti II | 0.04 | 0.08 | 0.10 | 0.03 | 0.14 |
V I | 0.17 | 0.00 | 0.06 | 0.01 | 0.18 |
Cr I | 0.11 | 0.00 | 0.06 | 0.02 | 0.13 |
Cr II | 0.02 | 0.07 | 0.09 | 0.05 | 0.13 |
Mn I | 0.12 | 0.00 | 0.08 | 0.03 | 0.15 |
Fe I | 0.08 | 0.02 | 0.11 | 0.06 | 0.15 |
Fe II | 0.00 | 0.07 | 0.10 | 0.05 | 0.13 |
Co I | 0.17 | 0.00 | 0.06 | 0.02 | 0.18 |
Ni I | 0.11 | 0.00 | 0.09 | 0.03 | 0.15 |
Ce II | 0.07 | 0.08 | 0.04 | 0.01 | 0.11 |
Nd II | 0.09 | 0.08 | 0.06 | 0.02 | 0.14 |
Sm II | 0.07 | 0.07 | 0.03 | 0.02 | 0.11 |
Notes: The effective temperature (
![]() ![]() the microturbulence ( ![]() ![]() is given as the sqaure root of the square sum of the individual errors. |
Acknowledgements
This research was supported in part by the Research Fund of the University of Istanbul, project numbers 1449/05052000 and OR-145/06112000. SJA thanks Dr. James E. Hesser, Director of the Dominion Astrophysical Observatory for the observing time. His contribution to this paper was supported in part by grants from The Citadel Foundation.