PASJ: Publ. Astron. Soc. Japan 55, 1133–1139, 2003 December 25
c 2003. Astronomical Society of Japan.
Abundance Analysis of the Silicon Star HR 6958
Ken-ichi K ATO
Osaka Science Museum, Nakanoshima, Kita-ku, Osaka 530-0005
kato@sci-museum.kita.osaka.jp(Received 2002 September 29; accepted 2003 August 27)

Abstract
The elemental composition of the chemically peculiar star HR 6958 has been studied with emphasis on
doubly ionized rare earths. A visual region spectrum taken with the ELODIE spectrograph at the Haute-Provence
Observatory was analyzed. A total of 40 elements including Pr III, Nd III, Tb III, Dy III, Ho III, and Er III were identified and their abundances computed. He is deficient by over 1 dex with respect to the Sun; the light elements (C–Ca),except for Si, have solar abundances; the iron group elements (Sc–Fe) are overabundant by 1 dex to 2 dex, with Ti
and Cr highly overabundant; and the lanthanide rare earths are overabundant by 3 dex to 4 dex. This abundance
pattern with He deficient, Si, Ti, Cr, Sr, and Pr overabundant indicates that HR 6958 is a member of Si stars.
Key words: stars: abundances — stars: chemically peculiar — stars: individual (HR 6958)

1. Introduction
Since the classification of stars with peculiar spectralfeatures into an apparent temperature sequence was made by
Morgan (1933), a large number of studies have been made toclarify the nature of a group of these stars. Preston (1974)classified the peculiar stars into four groups. The magnetic
chemically peculiar (CP) stars are designated as CP2, which
consist of two groups having different effective temperatures: acooler one, which shows an enhancement of elements, such as
Cr, Sr, and Eu; and a warmer group, which exhibits an unusual
strength of Si lines.
HR 6958 (= HD 170973, MV Ser) is an A0pSiCr CP2 starand its brightness changes by 0.03 mag (in the U band) witha period of 0.945099 d as cited in the Bright Star Catalogue(Hoffleit, Warren 1991). Recently, Adelman (1997) has revised
this period to 18.065 d based on his own uvby photometry.
López-Garcı́a et al. (2001) performed a fine analysis of HR
6958 for 27 atomic species. He I lines are not seen. The lightelements are mostly of solar abundance, except for Si, which isoverabundant by a factor 4. All of the heavy elements except
for Ni, which is solar, are greatly overabundant. In addition,
C is underabundant, S and Ca are overabundant, and the ironpeak elements are typically 10-times overabundant, except forTi and Cr, which are overabundant by factors of 75 and 35,
respectively. Mn is about 50-times solar, Sr is overabundant by
1000, and the rare earths are 1000 or more times overabundant.
Cowley (1984) emphasized that the behavior of lanthanide
rare earths provides a key to understand the surface chemistryof CP2 stars. In a previous study, Cowley (1976) tried to
identify rare earths in HR 6958 by using a statistical methodof wavelength coincidences, and showed the possible presenceof singly ionized lines for Ce, Pr, Nd, Sm, Eu, Gd, Tb, and Dy.
In later work, Mathys and Cowley (1992) reported the identification of Pr III lines in magnetic CP stars, including HR 6958.
With the efforts of atomic physicists, some data on oscillator
strengths are now available for doubly ionized rare earths (e.g.,
Bord et al. 1997; Biémont et al. 2002). Based on these data, theabundances or presence of doubly ionized rare earths have beenstudied for several CP stars (e.g., Cowley et al. 2000; Wahlgren
et al. 2000; Dolk et al. 2002; Nishimura et al. 2003).
2.

Spectrogram and Measurement of Lines

We used an echelle spectrum of HR 6958 obtained with
the ELODIE spectrograph at the Haute-Provence Observatory,
which was chosen from the HYPERCAT database (Prugniel,
Soubiran 2001). The spectrum, ranging from 4100 Å to 6800 Å
with a nominal resolution of 42000, is supplied after the spikes
due to cosmic rays, and the features of telluric lines were
removed from the crude spectrogram. The sequential numberof the spectrum used in this study was 00602.fits, the observation being made on 1999 July 22, with an exposure time of
40 min. Its S/N ratio was 84.8 at 5550 Å.
In addition, a high-dispersion photographic 2.4 Å mm−1
spectrum in the violet–blue region (λλ 3800–4600 Å), obtainedat the Dominion Astrophysical Observatory (DAO), wasreferred to in order to measure Hγ line profiles and also inorder to compare the equivalent widths of other lines withthose of ELODIE. The plate number was 10010. Its observational date was 1975 August 13. The spectrogram digitizedusing a PDS machine during the stay of Professor K. Sadakaneof Osaka Kyoiku University at the DAO in 1990 was kindly
provided to the author.
We computed the synthesized spectra by adopting the sameatmospheric parameters as López-Garcı́a et al. (2001) to makeline identifications and to select clean or minimally contaminated lines. We used the program SPTOOL, modified anddeveloped by Dr. Y. Takeda from the Kurucz’s programs
SYNTHE and WIDTH (Kurucz 1993a,b). The synthesized
spectra were also used to draw the continuum line.
The equivalent widths of lines were measured by assuming
a Gaussian profile with the program Nijiboshi written by
T. Hasui. We identified 578 lines as being suitable forabundance computations, and an additional 303 unidentified lines were also measured, whose equivalent widths were1134

K. KatoFig. 1. Comparison of the equivalent widths (in mÅ) between thisstudy and López-Garcı́a et al. (2001) for HR 6958.

between 9 mÅ to 126 mÅ. These identification lists are available at the web site of Osaka Science Museum.1
The measured lines were compared with the results ofLópez-Garcı́a et al. (2001). They used four photographic
4.3 Å mm−1 spectrograms of HR 6958, obtained with the
2.5 m telescope at the Mt. Wilson Observatory. Their equivalent widths are systematically lower than those of this study.
In figure 1, we compare the equivalent width in this study(ordinate) with that of López-Garcı́a et al. (2001) (abscissa).
A least-squares comparison of the equivalent widths Wλ ’s wasmade for 58 lines with equivalent widths of 10 mÅ to 120 mÅ:Wλ (this study) = (1.30 ± 0.08)Wλ (López-Garcı́a et al.)+ (1.36 ± 4.24).
(1)
Similarly, we compared the ELODIE data with equivalent
widths deduced from the DAO spectrogram. The result isWλ (this study) = (0.92 ± 0.05)Wλ (DAO)− (4.38 ± 3.42).

Atmospheric Parameters and the Magnetic Field

The effective temperature, Teff = 10750 K, and the surface
gravity, log g = 3.5, were obtained for HR 6958 by Adelman
and Rayle (2000) from spectrophotometric data and Hγ line1

http://www.sci-museum.kita.osaka.jp/˜kato/.

profiles (Adelman 1983; Adelman et al. 1989). They pointed
out that the measured fluxes beyond the longward part of the
λ 5200 feature are systematically brighter than model fluxes.
From a measurement of the DAO plate, we can confirm theresult for Hγ line profiles of Adelman and Rayle (2000). In
this study, we adopted their effective temperature and surface
gravity to make the analysis consistent with López-Garcı́a et al.
(2001). This effective temperature is slightly higher than thatof Sokolov (1998) (Teff = 10050 K). The rotational velocity of
HR 6958 has been reported to be 15 km s−1 (Hoffleit, Warren1991) or 10 km s−1 (Abt, Morrell 1995). We determined
its rotational velocity by fitting the ELODIE spectrum withsynthetic spectra computed for various rotational velocities.
The best fit was achieved for a rotational velocity, v sin i, of
8 km s−1 , which was adopted in this study.
With these parameters, we estimated the abundances forsome elements to determine the microturbulent velocity. The
condition used is that the elemental abundances should not
be dependent on the equivalent widths. The best values
found were 0.0 km s−1 for Ti II (62 lines) and Cr II (56 lines),
0.9kms−1 for Fe I (35 lines), 0.4kms−1 for Fe II (276 lines). On
the other hand, López-Garcı́a et al. (2001) obtained the slightly
different values such as 1.0kms−1 for Ti II, 1.4kms−1 for Cr II,
0.6kms−1 for Fe I, and 0.0kms−1 for Fe II. Finally, we adopted
a microturbulent velocity of 0.5 km s−1 as a mean value in theabundance computation.
The effects of a magnetic field on the spectral lines were
investigated to compare the strengths of the Fe II lines at
4385.38 Å and 4416.81 Å. These lines were studied in detailby Takeda (1991) as a line pair suitable for searching stellar
magnetic fields. They have the same multiplet number, and thegf values are essentially the same, while their effective Landéfactors (Z) are different (Z = 1.333 and 0.833 for 4385.38 Å
and 4416.82 Å, respectively). The present result for HR 6958
shows both lines to be the same strength; their measured equivalent widths are 106.6 mÅ and 103.6 mÅ for 4385.38 Å and
4416.82 Å, respectively. Moreover, the Fe II line at 6149.24 Å
(Z = 1.333) is a single component. These features suggest thatthe effect of a magnetic field is not so remarkable on HR 6958,
as already pointed by Mathys and Lanz (1992).

(2)

The difference between the Mt. Wilson spectrum and CCDbased ELODIE spectrogram is up to 30%. This is partly due tothe difference in the S/N ratio, scattered light, accuracy of thecontinuum placement between the photographic spectrum and
CCD-based Echelle spectrograms. However, the difference
mentioned above seems to be somewhat large to be accounted
for in terms of these reasons. HR 6958 may be a spectrum
variable, like a spectroscopic binary, as suggested by Catalano
and Renson (1998).
3.

[Vol. 55,

4.

Chemical Abundances

We computed the chemical abundances from the equivalent
widths with WIDTH9 (Kurucz 1993a), supplemented by thespectrum synthesizing program SPTOOL to determine theirabundances from lines heavily blended or affected by neighboring lines. Ten-times the solar metal content ATLAS model
atmospheres (Kurucz 1993a) were employed in the computation.
No effect of Zeeman broadening was taken into accountbecause of the weak magnetic field on HR 6958. The oscillator
strengths were taken from DREAM database (e.g., Biémont
et al. 2002)2 and Kurucz and Bell (1995) for rare earths
and other elements, respectively. The hyperfine splitting wasapplied to three Eu II lines listed in table 1 by adopting datafrom Kurucz and Bell (1995).
2

see also, http://www.umh.ac.be/˜astro/dream.shtml.

No. 6]Abundances of the Silicon Star HR 6958

1135

Table 1. Line list of lanthanide rare earth elements.∗

Atom

λ
(Å)

χ
(eV)

log gf

Wλ †
(mÅ)

log ε
mean

Atom

λ
(Å)

χ
(eV)

log gf

Wλ †
(mÅ)

log ε
mean

La II

4238.390

0.403−0.28010

Nd III

Ce II

4137.680
4144.500
4152.000
4165.630
4185.400
4214.030
4248.680
4373.800
4382.180
4429.290
4460.230
4463.430
4483.940
4523.060
4527.350
4560.290
4560.900
4562.380
5274.260

0.517
0.478
0.684
0.910
0.417
0.609
0.684
0.561
0.684
1.088
0.478
0.957
0.864
0.517
0.320
0.910
0.684
0.478
1.044

0.440
−0.2300.510
0.530−0.560−0.4700.140−0.360
0.200
0.330
0.320
0.080
0.150
−0.030−0.1100.310−0.170
0.230
0.150

47
27
35
42
9
24
25
13
18
36
30
20
21
28
24
25
10
32
32

4483.400
4903.260
4927.490
5085.000
5633.541
5677.180
5845.000
6145.050
6327.260
6690.860

0.631
0.000
0.461
0.296
0.141
0.631
0.631
0.296
0.141
0.461

−1.340−1.830−0.830−0.660−2.060−1.410−1.130−1.290−1.360−2.310

48
syn
63
72
syn
49
57
57
55
29

Sm II

4279.74378.235

0.277
0.659

−0.798−0.590

syn
syn

Eu II

4129.725
4205.042
4435.578
6645.030

0.000
0.000
0.207
1.380

0.204
0.117
−0.092
0.204

syn
syn
syn
23

Gd II

4215.022

0.427−0.55020

Tb III

Pr II

4206.719
5173.902
5815.331

0.550
0.968
1.591

0.480
0.340
0.271

syn
syn
syn

5847.232
6092.896
6687.698

0.348
0.587
1.027

−0.980−1.110−1.33028
22
syn

4437.610
4612.010
4642.260
4725.590
4775.300
4929.140
4964.580
5299.990
5449.380
5765.300
5844.4505998.9506090.000
6160.240
6195.630
6361.660
6500.080
6501.540
6616.529
6706.750

0.000
1.759
0.960
2.078
1.947
0.359
0.173
0.359
0.359
1.549
1.244
0.173
0.359
0.173
0.000
0.173
1.722
1.458
1.549
0.552

−3.400−1.260−1.770−1.320−1.280−1.880−2.210−0.530−2.190−1.100−0.850−1.800−0.820−0.980−1.040−2.080−1.140−1.400−1.720−1.640Dy II

Pr III

12
28
18
28
23
24
25
80
30
22
53
44
60
61
63
20
27
36
16
32

4364.2074754.992

2.427
2.417

0.260
−0.090syn
syn

Dy III

4355.296
4363.489
4401.567
4409.897
4434.269
4502.917
4510.027
4995.538
5099.734

0.516
1.151
0.881
1.366
0.516
1.151
0.881
0.881
0.881

−1.610−1.720−1.430−1.620−1.270−1.950−1.850−2.250−2.440

syn
syn
33
syn
syn
24
syn
syn
syn

Ho III

4416.250
4494.523

0.000
0.000

−1.550−1.360

29
syn

Er II

4630.882

1.951

0.033

syn

Er III

4706.580
5130.600

0.000
1.304

−0.880
0.570

10
34

4.10
4.10
5.59
5.38
5.08
5.47
4.92
5.58
5.06
5.02
4.73
5.51
4.94
5.08
4.98
5.22
5.08
5.02
4.75
5.07
5.46
5.15
4.7
4.9
4.9
4.83
5.66
5.10
4.78
5.33
5.02
4.80
5.11
5.55
5.36
4.65
5.38
5.39
5.13
5.20
5.25
4.80
5.00
5.44
5.07
5.04
5.15
5.44
5.63
5.54

4356.549
4422.410
4540.700
4735.510
5068.447

1.639
0.000
0.000
0.630
1.546

−1.460−1.740−2.540−1.580−1.660

syn
24
16
syn
6

5.37
4.90
5.34
5.39
5.20
5.48
5.48
5.43
5.36
5.53
5.39
5.0
4.9
4.95
4.3
4.0
4.1
4.37
4.25
4.85
4.85
4.03
4.07
4.1
4.06
4.9
4.8
4.85
4.6
4.6
4.77
4.9
4.9
5.07
4.9
5.0
5.2
4.89
4.16
3.9
4.10
4.48
4.48
4.47
4.13
4.61
3.9
3.97
4.23Nd II∗ The abundance, log ε, is expressed in the scale of log ε(H) = 12.00.

† The symbol syn indicates that the abundance was derived using themethod of spectrum synthesis.

1136

K. Kato[Vol. 55,

Table 2. Chemical composition of HR 6958 in logarithmic units compared with the study of López-Garcı́a et al. (2001), the Sun, cool magnetic CPstars, and three Si stars.

Atom

n

This
study

Standard
deviation

He I
C II
OI
Na I
Mg I
Mg II
Al I
Al II
Si II
Si III
P II
S II
Ca II
Sc II
Ti II
V II
Cr I
Cr II
Mn II
Fe I
Fe II
Fe III
Sr II
Zr II
Ba II
La II
Ce II
Pr II
Pr III
Nd II
Nd III
Sm II
Eu II
Gd II
Tb III
Dy II
Dy III
Ho IIIEr II
Er III

≤3
1
2
1
1
5
1
5
9
2
3
1
1
2
62
3
2
56
9
35
276
4
3
2
1
1
19
3
20
2
10
2
4
1
3
2
9
2
1
5

≤ 9.99
8.60
9.09
7.23
7.00
7.30
6.47
≤ 7.40
8.53
8.92
5.65
7.33
6.66
3.69
7.46
5.10
7.59
7.48
6.62
8.94
9.14
9.25
5.84
4.57
≤ 2.40
4.10
5.15
4.83
5.15
5.54
5.39
4.95
4.25
4.85
4.06
4.85
4.89
4.10
4.48
4.23

···
···
0.01
···
···
0.25
···
···
0.16
0.08
0.50
···
···
0.11
0.27
0.20
0.47
0.26
0.19
0.21
0.27
0.41
0.31
0.03
···
···
0.27
0.12
0.27
0.13
0.18
0.07
0.17
···
0.04
0.07
0.20
0.18
···
0.31

Teff
log g

10750
3.50

López-Garcı́a
et al. (2001)
8.26

7.10
7.48
6.54
8.18
7.84
7.81
7.09
3.93
6.89

7.20
6.92
8.83
8.50
9.08
6.26
5.02

5.74
5.96
5.52
5.60
5.73
5.31
5.59

10750
3.5

The resultant abundances are exhibited in table 1, where
the lines of lanthanide rare earth elements employed in thisabundance work are listed. A full list of lines is available via
links from the URL appeared in the footnote 1. The final results
are presented in table 2. The second column shows the numberof lines used in the analysis. In the third and fourth columns,
the mean abundances and standard deviation for the meansare displayed for each element. The result of López-Garcı́a

Sun
10.99
8.55
8.87
6.33
7.58
7.58
6.47
6.47
7.55
7.55
5.45
7.33
6.36
3.17
5.02
4.00
5.67
5.67
5.39
7.50
7.50
7.50
2.97
2.60
2.13
1.17
1.58
0.71
0.71
1.50
1.50
1.01
0.51
1.12
−0.10
1.14
1.14
−0.26
0.93
0.93

Cool
magnetic CP

HR 2258
Si

HR 5597 HD 192913
Si
Si

8.22

8.92

8.84

7.48

6.99
7.28

7.74

6.90
7.30
6.67

7.62

8.12

8.64

7.94

6.54
6.88
2.89
6.19

7.16
6.28
5.98

7.33
6.67
3.32
6.40

7.85
7.00
8.68
8.72
8.82
4.99
4.47
3.32

8.40
7.08
7.24
8.74
8.64
9.32
6.97
4.78
3.48

5.14
5.31

4.86
5.02

5.14
5.34

3.92

4.66

5.28

5.19

3.89
4.30
4.48

3.85
4.82

3.56

5.66
6.71
4.56

6.82
3.19
4.70
3.92
6.67
6.98
8.21

4.95
4.17
2.20
2.23
3.36
3.09

7.31
6.98
6.45
8.52
8.34
8.30
6.37
4.19

5.44

11300
3.2

11200
3.84

10900
3.40

et al. (2001) for HR 6958 is exhibited in the fifth column. The
abundances of the Sun (Grevesse et al. 1996), cool magnetic
CP stars (Adelman 1973; Adelman, Cowley 1986), and three
Si stars as samples of moderately hot CP2 stars are also given
for comparison.

No. 6]Abundances of the Silicon Star HR 6958

1137

4.1. Helium
The neutral helium lines in HR 6958 are too weak to detect.
We failed to find a line at 6678.15 Å, which is well-known to bea strong He I line in usual B and A stars, although the spectral
region around the line is fairly clear. HR 6958 is helium-poor
relative to the Sun, with a value typical of Si stars. Thus, we
can confirm the result of López-Garcı́a et al. (2001) for the He
abundance.
4.2. Light Elements (C to Ca)
Lines of C II, O I, Na I, Mg I, Mg II, Al I, Al II, Si II, Si III,P II, S II, and Ca II lines were measured and their abundances,or upper limits of the abundances, were computed. However, asmall number of elements give reliable abundances. The final
abundances are nearly solar, except for Si, which is greater by
1 dex than the solar abundance.
4.3. Iron Group Elements (Sc to Fe)
Sc is slightly enhanced, and the other iron group elements
show overabundances by 1.1 dex to 2.4 dex. In particular, Ti is
about 300-times the solar value. For Fe, doubly ionized lines
were also detected as well as neutral and singly ionized ones. A
complete ionization balance for all three species of iron can beachieved, if we adopt a higher effective temperature model ofTeff = 11100K and logg = 3.5. Although we can find features of
Ni and Co lines from a comparison of the synthesized spectrawith the observation, all of them show no clean line profiles
suitable for an abundance computation.
4.4. Heavy Elements (Sr, Zr, and Ba)
These elements are enhanced by 0.3 dex to 2.8 dex. The
resonance line of Ba at 4554.03 Å is fairly weak and subject
to blending with the Zr II line at 4553.934 Å. No other Ba line
was detectable in the spectral range of this study.
4.5. Rare Earths
Lines were identified using sources, including the line listof Kurucz and Bell (1995), National Institute of Standard andTechnology (NIST)3 , and DREAM database (Biémont et al.
2002)
The lines of singly ionized rare earths were detected andmeasured for La, Ce, Pr, Nd, Sm, Eu, Gd, Dy, and Er. The
Lu II lines seem to be present, but this has not been confirmed.
Doubly ionized rare earths lines of Pr, Nd, Tb, Dy, Ho,and Er were measured and their abundances evaluated. In
addition, four Gd III lines listed in the NIST atomic spectra
database3 were possibly identified, but their abundances were
not obtained due to a lack of atomic data. We failed to detect
the Ce III lines listed in Bord et al. (1997). Tm III lines as well as
Yb III and Lu III lines in the DREAM database (Biémont et al.
2002) were not seen in the spectrum. It was impossible to findthe Ba III and Eu III lines in the NIST atomic spectra database.
Their elemental abundances are in table 1.
HR 6958 is very enhanced in rare earth elements by 2.8 dex
to 4.4 dex relative to the Sun. Pr III is highly overabundantcomparing with those of the neighboring atoms (Ce and Nd),3

http://physics.nist.gov/PhysRefData/contents-atomic.html.

Fig. 2. Abundances of HR 6958 for (a) light elements, (b) irongroup elements, and (c) Ba and rare earths, compared to the resultof López-Garcı́a et al. (2001) for HR 6958, the solar abundances(Grevesse et al. 1996), and the abundances of cool magnetic CP stars(Adelman 1973; Adelman, Cowley 1986) and the silicon star HD
192913 (López-Garcı́a, Adelman 1999).

although Eu is normally enhanced and holds the odd–even rule.
Thus, Eu is not selectively overabundant on HR 6958. This
seems to be specific to HR 6958 (generally, to silicon stars)
when compared with the mean abundances of cool magnetic
CP stars (Adelman 1973; Adelman, Cowley 1986).
5.

Results and Discussion

Table 2 and figure 2 present the abundances obtained in thisanalysis with those of López-Garcı́a et al. (2001) for HR 6958,

1138

K. Katotogether with abundances of the Sun, cool CP2 stars, and Si
stars.
When comparing the abundances in this study with thoseof López-Garcı́a et al. (2001), Ti II, Cr II, and Fe II (for whicha fairly large number of lines are available in the abundancecomputation) show clearly large discrepancies of more than
0.3 dex. This result reflects the disagreement concerning theequivalent widths noted in section 2. However, the abundancesof rare earths in this study are systematically lower than theirresults.
Compared with the Sun, light elements are mostly of solar
abundances, except for He, Na, and Si. An interesting featureof light elements is an underabundance of magnesium. Leoneet al. (1997) showed that the magnesium abundance tends tobe lower in peculiar stars, especially in silicon stars; magnesium is obviously less abundant than that of main-sequencestars (figures 1 and 2 in their paper). HR 6958 seems to bea typical star with a lower magnesium abundance. Ca and Sc
are subsolar, while other iron group elements are on the orderof 10-times overabundant, except for Ti, which has 300-times
solar abundance. Sr and Zr are on the order of 100-times solar
abundances. As shown in figure 2c, there is a clear gap in theabundance between Ba and rare earths. Ba is not so enhanced,
while rare earths show fairly large excess, and preserve theodd–even rule. These features mark HR 6958 as a Si star
(Jaschek, Jaschek 1971; Jaschek, Jaschek 1987).
When compared with the mean abundances of cool magnetic
CP stars, there are a few disagreements in the abundances for
Si, Ti, Fe, Sr, and Ce.
HR 5049 (Teff = 10300 K, log g = 3.81) is a B-type magnetic
CP star with a strong magnetic field of 4.7 kG. The atmosphericparameters are similar to those of HR 6958 (Teff = 10750 K,
log g = 3.5). Recently, Nishimura et al. (2003) have studied thechemical composition of HR 5049 based on the spectra taken at
the European Southern Observatory, and revealed the extreme
deficiency of He and the overabundances of Cl, Co, and rareearths. Since Cl and Co lines of HR 6958 are very weak, it was
impossible to detect their lines. On the other hand, Si and Ti of
HR 6958 are enhanced more than those of HR 5049.
Table 2 compares the final abundances of HR 6958
with those of three silicon stars: HR 2258 (= HD 43819;
Teff = 11300 K, log g = 3.20; López-Garcı́a, Adelman 1994),
HR 5597 (= HD 133029; Teff = 11200 K, log g = 3.84; LópezGarcı́a, Adelman 1999), and HD 192913 (Teff = 10900 K,[Vol. 55,

log g = 3.40; López-Garcı́a, Adelman 1999). The agreement
between HR 6968 with these stars is quite good. The exception is Ti II, whose abundance in our analysis is greater by
1 dex than that of the three silicon stars. The abundance patternof HR 6958 is very similar to that of HD 192913, except for
Eu, which is extremely enhanced by 6.20 dex compared to the
Sun in HD 192913. Since the model parameters of both stars
closely resemble each other, the abundance pattern supports
HR 6958 to be a silicon star.
Recently, Dolk et al. (2002) and Nishimura et al. (2003)
have shown that the abundance ratio between Nd and Pr in CP
stars (Hg-Mn stars, Am stars, and CP2 stars) decreases linearlywith the effective temperature. It is interesting to know thatthe logarithmic abundance ratio of + 0.24 dex (the abundance
difference between Nd III and Pr III) for HR 6958 is just on thesequence.
After determining the final abundances, we retried to synthesize the spectrum. In spite of such an effort, a total of 303
lines remained unidentified. The rate of the unidentified lines
with respect to the total measurable lines is about 20%, which
is fairly greater than the rate of 7% for the case of the coolmagnetic CP star HR 7575 (Kato, Sadakane 1999). One halfof the unidentified lines have candidates (most of them are linesof singly ionized rare earths and iron group elements), but are
too weak to reproduce the observed line features. Although
about 70% of the unidentified lines can be seen in the spectrumof the cool magnetic CP star 10 Aql (Teff = 7750 K, log g =4.0; Ryabchikova et al. 2000), only about 40% of them arepresented in the spectrum of the hot magnetic CP star α 2 CVn,
whose effective temperature (Teff = 10410 K; Sokolov 1998)
closely resembles that of HR 6958. This implies that most of
the unidentified lines can be attributed to some atoms abundantboth in HR 6958 and in 10 Aql but less in α 2 CVn than in
HR 6958. To identify these unknown elements and lines, moreaccurate atomic line data are desirable.
We thank the Haute-Provence Observatory for kindly
providing us the data on HR 6958, and Dr. K. Sadakane forthe use of the photographic spectral data of the DAO. We
also appreciate the assistance of Dr. Y. Takeda and T. Hasui
for making their program codes available to the author,
and acknowledge helpful comments from the referee Dr. S.
Adelman.

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