A&A 374, 957-967 (2001)
DOI: 10.1051/0004-6361:20010645
E. Solano1 - E. Paunzen2 -
O. I. Pintado3, - J. Varela4
1 - Laboratorio de Astrofísica Espacial y Física Fundamental
(LAEFF), Apartado de Correos 50727, 28080 Madrid, Spain
2 -
Institut für Astronomie der Universität Wien, Türkenschanzstr. 17,
1180 Wien,
Austria
3 -
Departamento de Física, Facultad de Ciencias Exactas y Tecnología,
Universidad Nacional de Tucumán, Argentina - Consejo Nacional de
Investigaciones Científicas y Técnicas de la República
Argentina
4 -
IMAFF, Consejo Superior de Investigaciones Científicas, c/ Serrano 113,
28006 Madrid, Spain
Received 7 December 2000 / Accepted 26 April 2001
Abstract
This is the first of two papers whose main goal is to update and
improve the
information available on the physical
properties of the Bootis stars. The determination of the stellar
parameters is of fundamental importance to shed light into the different
theories proposed to explain the
Bootis phenomenon.
With this aim, projected rotational velocities,
effective temperatures, surface gravities and chemical abundances of a sample
of suspected
Bootis stars have been calculated. Five objects showing
composite spectra typical of binary systems were found in our analysis. The
abundance distribution of the program stars does not resemble the chemical
composition of the class prototype,
Boo, which poses some concerns regarding
the idea of a well-defined, chemically homogeneous group of stars. A
possible relation between rotational
velocities and the
Bootis phenomenon has been found. This result would be
in agreement with the accretion scenario proposed by
Turcotte & Charbonneau (1993).
Key words: stars: chemically peculiar - stars: abundances
Observational results supporting the accretion theory are the high [C/Si]
and [O/Si] abundance ratios found in some Bootis stars which can be used as
indicators of gas-dust separation (Paunzen et al. 1999a): carbon
and oxygen
have low
condensation temperatures and tends to
remain in the gas phase of the interstellar medium. Silicon, on the other
side, has a high condensation temperature representing
the elements locked up in the grains. Any preferential accretion of gas will
lead to a [C/Si] or [O/Si] larger than solar. The
correlation with Si of other elements of high condensation temperature also
fits nicely with the accretion scenario (Stürenburg 1993).
Charbonneau (1991) combined the accretion hypothesis with the
diffusion
theory concluding that, with an accretion rate of
the order of 10
yr-1, many peculiar
characteristics
of
the
Bootis stars (including their restriction to the above-mentioned spectral types)
are
reproduced quite naturally. Turcotte & Charbonneau (1993)
investigated the
effect of mixing through meridional
circulation, concluding that the abundance anomalies would disappear in
106 years, so all of the
Bootis stars should be
essentially young objects on or just arriving to the ZAMS. The scarcity of
Bootis will be also explained by the
strict requirements for the accretion rate. The fact that all
Bootis stars are
young objects is further strengthened by the lack of
Bootis stars in open
clusters older than 107 years (Gray & Corbally 1988) and the
discovery of
Bootis
stars in the young Orion OB1 association and NGC 2264 (Paunzen & Gray
1997).
A critical point in the accretion scenario is the necessary existence
of gas and dust
shells around Bootis stars. Some
Bootis stars show clear indications
of both gas (circumstellar lines) and
dust (IR excess) shells (Holweger & Rentzsch-Holm 1995;
Holweger et al. 1999).
However, the fact that some
Bootis stars do not show
evidence of shells does not necessary rule out this hypothesis. King
(1994) concluded that the amount of depleted gas required
to cause underabundances in
Bootis stars is small enough that any circumstellar
dust associated with this gas is not necessarily
detectable in the IR or submillimetre regions.
Michaud & Charland (1986) investigated the effect on abundances
in stellar atmospheres of the diffusion mechanism operating in the
presence of mass
loss. They found that, after about 108-109 years,
mass-loss
rates of only 10
yr-1 can reduce the extreme
overabundances predicted by diffusion theory for the Am
stars to underabundances of many elements, the degree
of underabundance of an element being a function of both gravity and time.
According to this,
Bootis stars would be rather old and
at the end of their Main Sequence.
There are, however, some points that cannot be explained with theory. One is
the moderate underabundances predicted (only a factor of
five less solar), very far from the strong underabundances found by
Venn & Lambert (1990) in the
three classical Bootis stars (
Boo,
Ori and 29 Cyg). Also it is
difficult to understand how
diffusion could operate in the presence of the meridional circulation which
would likely be generated by rapid rotation. Turcotte & Charbonneau
(1993) showed that
even an equatorial
rotational velocity of
50 kms-1 suppresses the appearance, at any epoch of Main
Sequence evolution, of the characteristic
Bootis abundance pattern.
Faraggiana & Bonifacio (1999) proposed this
alternative to explain the underabundances of, at least,
some Bootis stars. The composite spectra of a binary system with disentangled
components not very dissimilar produces a veiling
effect with apparent underabundances. The authors pointed out that the lack
of a uniform pattern in the chemical composition of the
Bootis stars can be
easily reconciled with the binary hypothesis. Moreover, they
claimed that the hypothesis of all
Bootis stars being very young objects in the
late phase of their PMS evolution is highly improbable on the basis of the
rapid evolution and the number of bright
Bootis candidates which would imply
that the
star formation process is
still very active in the solar neighborhood, leading to an unexpected large
number of Main Sequence B stars in a similar volume.
Andrievsky (1997) proposed a complementary scenario in which Bootis stars would be the result of the coalescence of contact binaries of W UMa type.
The system would be formed by two main-sequence components of approximately
equal spectral types. This scenario would lead to an age for
Bootis stars of
1 Gyr explaining the rather evolved nature of some
Bootis stars and the origin
of the material around them.
A class definition of the Bootis group is essential
on one hand to distinguish these stars from other groups of stars
populating the same region of the HR diagram and, on the other hand, to shed
light into the nature of the
Bootis phenomenon. Ideally, a stellar class should be
formed by a homogeneous sample of stars showing common properties originated
by the same astrophysical processes. Historically, this has not been the case
for the
Bootis group since the use in the past of classification criteria not
unique to the group (weakness of the Mg II 4481 Å line,
presence of spectral features at 1600 Å and 3040 Å, IR excess, ...)
has led to the inclusion of spurious members (horizontal branch stars,
Ap stars, shell stars, He-weak stars, ...). Although the problem of the class
definition can be efficiently alleviated making use of
unambiguous criteria defined in the ultraviolet range
(Solano & Paunzen 1998,
1999), the chemical composition-based definition of the class
makes it necessary to perform an accurate determination of the stellar
parameters for the final decision on the membership of a potential candidate
to the
Bootis group. The
observed sample was selected from the list of suspected
Bootis stars given in
Paunzen (2000). HD 68758 and HD 184190 were also included as
possible members
of the class.
The observations were performed in May 1998 at the Complejo
Astronómico el Leoncito (CASLEO) using the 2.15-m telescope equipped with
a REOSC
echelle spectrograph and a Tek-1024 CCD. A grating with 1200 lines mm-1 was
used as a cross disperser. A resolving power of 26000 was achieved. Two
wavelength ranges were observed (3830-4570 Å; 4520-5230 Å).
The spectra were reduced using the context echelle of the MIDAS reduction package. A typical session of echelle reduction comprises the following steps: bias subtraction, spatial positioning of the spectral orders, flatfield correction (several master flatfield exposures were taken every night), background subtraction, order extraction and wavelength calibration. The wavelength calibration of the stellar spectra was done with thorium-argon comparison spectra. Polynomial calibrations of wavelength as a function of pixel number were calculated for every night. The lower sensitivity at the edges of the orders typical of echelle spectra (ripple effect) was corrected by using Procyon as standard star. A blaze function was derived for each order.
The reliability of the measured equivalent
widths depends to a great extent on the accuracy of the continuum
placement: an improper placement of the continuum level will lead to
systematic errors which
can misrepresent the data. Most Bootis stars are characterized as having broad and
often shallow
absorption line profiles. In this case, the continuum placement cannot be set
by simply
connecting the highest points in the observed spectrum, which would produce an
underestimation of the
equivalent widths. We solved this
problem by defining the "true'' continuum level as that of a
synthetic spectrum of physical parameters (
,
,
[M/H],
)
similar to those of the observed object. As a first approach, effective
temperatures and surface gravities were
derived using Moon & Dworetsky (1985). Rotational velocities were
estimated using the method described in
Sect. 3 where a pseudo-continuum was defined by connecting the highest
points in the line profile: the
small influence
of the continuum level on the method used to calculate rotational velocities
permits this
approximation. Also, a value of
was used for all the
stars, which is justified by the results from our detailed abundance analysis.
The normalized spectra were
then derived by dividing the extracted spectra by this continuum level. Line
identification was performed with the help of Moore et al. (1966).
To have an estimation of the measurements errors in equivalent widths, the
standard star
Procyon
(
CMi, HR 2943) was observed and the spectra compared to the spectrum
of the Atlas of Procyon (Griffin & Griffin 1979). It can be seen
in Table 1 how the equivalent widths
measured in the observed spectra are only
slightly larger than those measured in the Atlas. We also compared the
equivalent
widths of spectral lines present in the overlapping region between spectral
orders finding no systematic trend. The list of observed objects is
given in Table 2.
Observing | Relative | Number |
date | error | of lines |
08/05/1998 |
![]() |
12 |
09/05/1998 |
![]() |
9 |
10/05/1998 |
![]() |
11 |
11/05/1998 |
![]() |
11 |
12/05/1998 |
![]() |
20 |
13/05/1998 |
![]() |
17 |
Identification | Spectral. | V mag. | Region |
![]() |
![]() |
![]() |
type | (Å) | This work | This work | This work | ||
/Other sources | /Other sources | /Other sources | ||||
HD 68758 | A1IVp | 6.52 | 3830-4570 |
![]() |
||
3830-4570 | 8100(H![]() ![]() |
![]() |
||||
3830-4570 | 8050(H![]() |
![]() |
||||
3830-4570 | 8066(H![]() |
![]() |
||||
4520-5230 | 285 (1) | |||||
/8300(4) | /3.7(![]() |
|||||
8150 | 3.7 | 299 | ||||
HD 75654 | hF0mA5V | 6.38 | 3830-4570 | 7256(H![]() ![]() |
![]() |
|
4520-5230 | ![]() |
|||||
/7275(4),7200(10) | /3.85(![]() |
/45(10) | ||||
7250 | 3.8 | 44 | ||||
HD 81290 | kA5hF3mA5V | 8.89 | 3830-4570 | 6775(H![]() ![]() |
![]() |
|
3830-4570 |
![]() |
|||||
4520-5230 | 6800(H![]() |
![]() |
||||
4520-5230 | 6750(H![]() |
![]() |
||||
/6750(4),6760(6) | /3.5(4) | |||||
6800 | 3.5 | 56 | ||||
HD 83041 | kA2hF2mA2V | 8.80 | 3830-4570 |
![]() |
||
3830-4570 |
![]() |
|||||
3830-4570 | 6800(H![]() |
![]() |
||||
4520-5230 | 7050(H![]() |
![]() |
||||
4520-5230 | 6975(H![]() |
![]() |
||||
4520-5230 | 6843(H![]() |
![]() |
||||
/6850(4),6900(6) | /3.6(4),3.3(6) | |||||
6900 | 3.6 | 95 | ||||
HD 107233 | kA1hF0mA1Va A1V | 7.37 | 3830-4570 | 7000(H![]() ![]() |
![]() |
|
3830-4570 | 6900(H![]() ![]() |
![]() |
||||
/7225(4),7244(5),7200(2) | /4.2(![]() |
|||||
7000 | 4.1 | 111 | ||||
HD 109738 | kA1hA9mA1V | 8.29 | 3830-4570 | 7437(H![]() ![]() |
![]() |
|
3830-4570 | 7425(H![]() |
![]() |
||||
7500(H![]() |
![]() |
|||||
7500(H![]() |
![]() |
|||||
/7575(4),7603(6) | /3.9(4),3.8(6) | |||||
7450 | 3.9 | 166 | ||||
111005 | hF0mA3V | 7.97 | 3830-4570 | 7250(H![]() ![]() |
![]() |
|
3830-4570 | 7415(H![]() ![]() |
![]() |
||||
3830-4570 | 7525(H![]() ![]() |
![]() |
||||
4520-5230 | 7500(H![]() |
![]() |
||||
4520-5230 | 7500(H![]() |
![]() |
||||
/ 3.8(![]() |
||||||
7400 | 3.8 | 138 | ||||
HD 142703 | kA1hF0mA1Va | 6.13 | 3830-4570 | 7172(H![]() ![]() |
![]() |
|
4520-5230 |
![]() |
|||||
/7200(4), 7400(2),7294(6) | /4.1(![]() |
/100(10),95(12) | ||||
7100 | 4.0 | 117 | ||||
HD 142994 | A3Va | 7.18 | 3830-4570 | 6900(H![]() ![]() |
![]() |
|
3830-4570 | 7035(H![]() ![]() |
![]() |
||||
4520-5230 |
![]() |
|||||
4520-5230 |
![]() |
|||||
/7079(1),7244(2),7000(4) | /3.5(1,2,3),3.4(4) | /220(1),195(3) | ||||
6950 | 3.4 | 206 |
Identification | Spectral. | V mag. | Region |
![]() |
![]() |
![]() |
type | (Å) | This work | This work | This work | ||
/Other sources | /Other sources | /Other sources | ||||
HD 156954 | hF1mA5V | 7.69 | 3830-4570 | 6775(H![]() ![]() |
![]() |
|
3830-4570 | 7000(H![]() |
![]() |
||||
4520-5230 | 6937(H![]() |
![]() |
||||
/7050(4),7079(6) | /4.2(![]() |
|||||
7000 | 4.1 | 51 | ||||
HD 168740 | hA7mA2V | 6.13 | 3830-4570 | 7875(H![]() ![]() |
![]() |
|
4520-5230 | 7575(H![]() |
![]() |
||||
/7650(4),7700(10) | /4.15(![]() |
/147 (10) | ||||
7700 | 4.0 | 162 | ||||
HD 184190 | A8 | 9.74 | 4520-5230 | 7200(H![]() |
![]() |
|
7250(4) | 4.0(4) | |||||
7250 | 4.0 | 19 | ||||
HD 193281 | hA3mA2Vb | 6.30 | 3830-4570 | 8000(H![]() ![]() |
![]() |
|
4520-5230 |
![]() |
|||||
4530-5230 |
![]() |
|||||
/8100(4), 8100(10), 8080(13) | /3.5(4),3.6(10),3.6(13) | /95(10), 83(13) | ||||
8050 K | 3.5 | 103 | ||||
HD 204041 | A1Vb | 6.46 | 3830-4570 | 7857(H![]() ![]() |
![]() |
|
3830-4570 | 7950(H![]() ![]() |
![]() |
||||
4520-5230 | 7950(H![]() |
![]() |
||||
/8100(4),8128(6),8100(7) | /4.0(![]() |
/65(7),68(8),70(11) | ||||
8000 | 4.0 | 69 | ||||
HD 210111 | kA2hA7mA2Vas | 3830-4570 | ![]() |
|||
3830-4570 | ![]() |
|||||
4520-5230 | 7415(H![]() |
![]() |
||||
/7762(2),7603(5),7450(4) | /3.9(![]() |
/55(8),60(11) | ||||
7450 | 3.8 | 57 |
References:
(1) Paunzen et al. (1998a); (2) Iliev & Barzova (1995);
(3) Bohlender et al. (1999);
(4) Moon & Dworetsky (1985);
(5) Paunzen (1997);
(6) Paunzen et al. (1998b);
(7) Stürenburg (1993);
(8) Holweger & Rentzsch-Holm (1995);
(9) North et al. (1994);
(10) Paunzen et al. (1999a); (11) Gray & Corbally (1993);
(12) Faraggiana & Bonifacio (1999); (13) Holweger et al. (1999).
Five targets of our sample show a composite spectra and no attempt to derive their abundance pattern was done (Fig. 1). All of them are identified as members of binary systems:
![]() |
Figure 1: Observed stars with composite spectra. |
Open with DEXTER |
Different methods can be found in the literature to calculate
rotational velocities. In this work we have used the technique
based on the Fourier transform proposed by Gray (1992) which has shown to give
superior results than those based on the
convolution of a non-rotating standard star of similar spectral type
with a rotational function or on the
identification of a particular parameter of the spectral lines (e.g. FWHM).
Whereas these two latter methods require one to build up a
calibration of rotational velocities, Gray's method provides a direct and
independent measurement of .
In short, the method relies on
the relation between
and the frequencies where the
Fourier transform of the rotational profile
reaches a relative minimum. Instrumental
broadening may add relative minima in the
Fourier transform but at higher frequencies than those used to calculate
the rotational velocity. Projected rotational
velocities for our sample of stars are given in Table 2.
It is well established that the Balmer lines are good temperature
indicators for
8000-8250
because of their small
gravity
and metallicity dependence (Solano & Fernley 1997). As the spectral orders are not
wide enough to fully embrace the wings of the observed Balmer lines
(H
,
H
,
H
), an alternative method
based on the variations in intensity at two
different wavelengths was used instead.
These wavelengths, free of contamination of metallic lines,
were selected in such a way that they lie in the region of the Balmer line
profile where the dependency with temperature is maximum. Prior to this,
the observed Balmer lines were shifted to the laboratory values to
correct for radial velocity displacements. Effective
temperatures were then calculated by comparing the observed intensity ratio
with the intensity ratio measured in a grid of synthetic Balmer profiles
previously convolved with the corresponding rotational and instrumental
profiles.
Balmer lines are known to be very sensitive to convection as it can
alter the temperature structure where the lines are formed. In this work, the
turbulent convection model developed by Canuto et al. (1996) and
implemented in the ATLAS9 code by Kupka (1999) has been
used. This convection model has been found to
reproduce adequately the temperature of standard stars whereas the standard
mixing-length theory models with overshooting, originally implemented in ATLAS9,
are clearly discrepant
(Gardiner et al. 1999; Smalley & Kupka 1997). The effective temperatures are given
in Table 2. An error of
has been
assumed.
Some concern may also exist in the use of Balmer lines as temperature
indicators due to the existence of peculiar Balmer
profiles (weak, narrow core typical of late A-type dwarfs but with strong
wings proper to early A-type dwarfs) in some
Bootis stars (Gray 1988). Based on this criterion,
Bootis stars have been
traditionally classified in two groups:
normal hydrogen-lines (NHL) and peculiar hydrogen-line profiles (PHL).
Iliev & Barzova (1993) suggested that PHL profiles could be fitted by using two
theoretical profiles with different effective temperatures, one for the core
and other for the wings. Faraggiana & Bonifacio (1999) pointed out that this duplicity
can be explained in terms of a binary system although
some of the stars catalogued as PHL objects (e.g., HD 142703, HD 142994,
HD 204041) do not show in their spectra any
sign of the presence of companion. Moreover, this scenario would
not explain why PHL profiles are more frequent at
lower temperatures as proposed by Iliev & Barzova (1993).
To ensure the reliability of the temperatures calculated using Balmer lines,
they have been compared with those calculated
using the photometric calibration by Moon & Dworetsky (1985) finding no systematic
differences (Table 2).
Identification | Parallax | Bolometric |
![]() |
log
![]() |
log (
![]() |
log t | |
Correction | PMS | MS | |||||
HD 68758 |
![]() |
-0.32 | 0.37 | 1.74 | 0.38 | 6.21 | 8.63 |
HD 75654 |
![]() |
-0.15 | 1.77 | 1.18 | 0.27 | 6.92 | 8.97 |
HD 107233 |
![]() |
-0.17 | 2.65 | 0.82 | 0.20 | 7.17 | 8.93 |
HD 111005 |
![]() |
-0.15 | 1.62 | 1.23 | 0.28 | 6.84 | 8.61 |
HD 142703 |
![]() |
-0.25 | 2.26 | 0.98 | 0.22 | 7.00 | 9.13 |
HD 156954 |
![]() |
-0.09 | 3.03 | 0.67 | 0.17 | 7.30 | 8.50 |
HD 168740 |
![]() |
-0.20 | 1.66 | 1.22 | 0.23 | 6.90 | 9.32 |
HD 204041 |
![]() |
-0.32 | 1.43 | 1.31 | 0.29 | 6.90 | 8.91 |
HD 210111 |
![]() |
-0.25 | 1.65 | 1.22 | 0.28 | 6.80 | 8.31 |
The method of calculating spectroscopic gravities using the standard
technique of making
abundances of
both neutral and ionized iron lines agree was not used here due to the
scarcity of usable Fe II lines in the observed
wavelength range. Alternatively, surface gravities were derived for those
stars with accurate parallaxes in the HIPPARCOS catalogue
following a method similar to that described in Nissen et al. (1997). In short,
the method relies on the basic relations
![]() |
(1) |
![]() |
(2) |
The mass of the stars were calculated from its position in the
-log
diagram by interpolating in the isochrones given in
Claret (1995). The mass so calculated was used as input value in
the
log
-
diagram to obtain the surface gravity. A chemical
composition X = 0.70, Z = 0.02 (solar abundance) was adopted. This is based on the fact that
the main contribution to the overall metallicity is due to
C, N, O (solar abundant in
Bootis stars). Furthermore, there are strong indications that the
Bootis phenomenon is restricted to the stellar surface (Holweger &
Rentzsch-Holm 1995). The mixing length and core
overshooting parameters were fixed to 1.52 and 2.0 respectively. Surfaces
gravities are given in Table 2. Parallaxes, bolometric corrections,
magnitudes,
luminosities, masses and ages are displayed in Table 3. Age estimations for the sample
stars indicate that they cover an area slightly above the Main Sequence. Surface gravities calculated
using HIPPARCOS are only slightly higher
than those derived using Moon & Dworetsky (1985) (
). A conservative
value of
= 0.15 dex has been estimated based on errors in the
bolometric correction and in the calculated mass.
![]() |
Figure 2: Comparison between observed and synthetic spectra for two stars of our sample. |
Open with DEXTER |
The abundance analysis was performed making use of a modified version of the ATLAS9 code (Kurucz 1993) in which the classical mixing length theory used for the treatment of convection has been replaced by the turbulent theory implemented by Kupka (1999) as described in Sect. 4.
Since the wavelength coverage of our spectra is quite extensive, we have
been restrictive
in the selection of lines for abundance analysis. Weak lines with large errors
in equivalent width due to noise and strong lines with a high sensitivity to
errors in microturbulence were discarded. Moreover, the typically high
rotational
velocities of the Bootis stars produce a strong blending of the spectral lines
and it is an additional limiting factor in
the line selection. Only blended features formed by lines of the same element
and ionization stage were considered (Fig. 2).
The abundance of the different elements were computed line by line
by linearly interpolating in the grid of metallicity values for a given
and
,
the final abundance being the average value. Results are
displayed in Table 4 and plotted in form of abundance patterns
in Fig. 3. For those program stars in common with Stürenburg
(1993), abundances have been also displayed, showing good
agreement.
Microturbulence was calculated in an independent way for each individual spectrum by
making abundance results of weak and strong lines agree. We have used the
iron lines since they are the most numerous and are spread over a
wide range in equivalent width. In those cases where the scarcity of Fe I
lines prevents from using this method, a value of kms-1 was adopted (Stürenburg 1993).
One source of uncertainty is the quality of the oscillator strengths
values. To avoid using lines with
values of poor quality, a careful
selection has been made using the most recent
papers in the literature and the VALD database (Kupka et al. 1999). An average value was adopted.
As a further test we compared, for every spectrum, the abundance value
derived from every single line with the average value. Lines with abundance
values clearly
discrepant (which could be attributed to a wrong
)
were discarded.
Another source of uncertainty results from the errors of the input
atmospheric parameters. To get an estimation of how
sensitive our results are to these errors we have derived the abundances of
a synthetic spectrum with
:
7500
,
:
4.0, [M/H]: -1.0 varied by
,
dex. Errors in
metallicity of
dex are obtained on average.
As expected, the abundance
of ions changes with changes in
whereas the neutral species tend to
have
negligible changes. On the contrary, changes in
produce
changes in the
abundances derived from neutral species while the
abundances derived from ions remain unchanged.
Uncertainties in the abundance values due to NLTE effects should also be included. However, according to Heiter et al. (1998), corrections due to non-LTE effects are of the same order as the abundance error bars in the range of temperature of our program stars making the LTE approach perfectly valid.
Finally, it must be stressed that some abundances rely in only one or few spectral lines so that the uncertainty is high.
Identification |
![]() |
[MgI/H] | [CaI/H] | [ScII/H] | [TiII/H] | [MnI/H] |
[CrI/H] | [CrII/H] | [FeI/H] | [FeII/H] | |||
Procyon | 2.0 |
![]() |
![]() |
-0.02 (1/1) |
![]() |
-0.03 (1/1) |
![]() |
![]() |
![]() |
![]() |
|||
HD 68758 | 3.0 | |||||
-0.58 (1/1) | ||||||
HD 75654 | 3.5 |
![]() |
![]() |
![]() |
![]() |
|
![]() |
-1.07 (1/1) |
![]() |
![]() |
|||
HD 81290 | 3.0 |
![]() |
-0.89 (1/1) | -0.88 (1/1) |
![]() |
![]() |
-1.21 (1/1) |
![]() |
![]() |
||||
HD 83041 | 3.0 |
![]() |
-1.31 (1/1) | -0.86 (1/1) |
![]() |
|
![]() |
-1.33 (1/1) | |||||
HD 107233 | 4.0 |
![]() |
||||
-0.96 (1/1) |
![]() |
![]() |
||||
HD 109738 | 3.0 | |||||
![]() |
-1.25 (1/1) | |||||
HD 111005 | 3.0 |
![]() |
||||
![]() |
![]() |
|||||
HD 142703 | 3.0 |
![]() |
![]() |
|||
![]() |
![]() |
|||||
HD 142994 | 3.0 | |||||
![]() |
||||||
HD 156954 | 2.5 |
![]() |
![]() |
|||
-0.76 (1/1) |
![]() |
![]() |
||||
HD 168740 | 3.0 |
![]() |
||||
![]() |
-0.73 (1/1) | |||||
HD 184190 | 3.0 | 0.20 (1/1) | -0.18 (1/1) |
![]() |
||
![]() |
0.0 (1/1) |
![]() |
||||
HD 193281 | 4.0 |
![]() |
![]() |
|||
![]() |
![]() |
|||||
HD 204041 | 3.0 |
![]() |
![]() |
|||
-1.08 (1/1) |
![]() |
![]() |
||||
HD 210111 | 3.0 |
![]() |
![]() |
![]() |
||
![]() |
-1.15 (1/1) |
![]() |
![]() |
![]() |
Figure 4: Relation between physical parameters. |
Open with DEXTER |
In order to search for correlations providing clues about the nature of Bootis stars, the abundances for our program stars have been compared with those of
Procyon, a primary standard with solar metallicity (Steffen 1985). Chemical abundances of
the class prototype,
Boo, were taken
from the literature (Paunzen et al. 1999b). As can be
deduced from Table 4, all the program stars but one (HD 184190)
show a clear metal deficiency when compared to Procyon. Moreover, it is quite
evident
that the candidate
Bootis stars here analyzed do not resemble the abundance
pattern of the
class prototype, in particular for some chemical species (magnesium, titanium
and iron). This result poses some concern with the idea of that
Bootis stars form a well separate, chemically homogeneous group of the
stars. On the contrary, this result reinforces the hypothesis proposed by
Stürenburg (1993) that
Bootis stars cover a continuous
sequence of underabundances from very metal week to solar metallicities.
At this point, it would be necessary to reopen the question of the Bootis class
definition. If the definition
by Baschek & Searle (1969) - "...
Bootis stars can be
defined as stars whose composition resembles that of
Boo
itself'' - is strictly adopted, then all the stars in our sample should
be rejected as potential candidates to the
Bootis group. On the contrary, if the
more conservative definition
given by Paunzen et al. (1997a) is considered - "...
Bootis stars are Population I,
metal-weak (except for C, N, O, S) stars'' -, all the stars (except
HD 184190) could be catalogued
as members of the
Bootis class provided that they show solar abundances for
carbon, nitrogen, oxygen and sulphur. It is of fundamental importance to
stress this point since the
Bootis phenomenon is not ascribed to an overall metal deficiency but to a
mechanism able to produce underabundances of the heavy elements contrasting
with the solar abundances of C, N, O and S. None of these species have been
analyzed in this paper due to the lack of spectral lines fulfilling the
requirements quoted in Sect. 6. NLTE abundances of nitrogen and sulphur for
the program stars will be part of a subsequent paper
(Kamp et al. 2001).
The observed distribution of the abundance values does not permit us to discriminate
between the theories proposed to explain the Bootis phenomenon. It could be explained in
terms of the diffusion/mass-loss hypothesis on the basis of
different
stages of the diffusion process. Under the assumption of the accretion
hypothesis, the observed differential deficiencies would reflect different scenarios in
the gas-dust decoupling ascribed to different disk properties. Moreover, the binarity
theories would be also supported.
The relation between metallicity and
,
,
has been plotted in Fig. 4. Whereas no obvious correlation between
and
with [Fe/H]
exists, there is evidence for a connection between projected rotational
velocities and metallicity in the sense that the metallicity is higher when
increases, although the correlation coefficient (
)
and the
small number of measurements (n = 14) do not permit to obtain any statistically significant conclusion. A similar result was found by Holweger & Rentzsch-Holm
(1995) using
calcium abundances derived from the Ca II K line. If, after the
determination of the C, N, O and S abundances the membership of the program
stars to the
Bootis group is finally confirmed, this result would nicely fit with
the accretion theory: for large
the meridional circulation
mixes material of solar composition from the stellar interior into the
convection zone so that any surface contamination due accretion of
circumstellar material should vanish.
Acknowledgements
The authors acknowledge use of the CCD and data acquisition system supported under U.S. National Science Foundation grant AST-90-15827 to R. M. Rich.
ª![]() |
Figure 3:
Chemical abundances of the program stars (filled rectangles). Abundances from Stürenburg (1993)
(triangles) and Paunzen et~al. (1999b) (empty rectangles) are also plotted for comparison. Abundances for the class prototype (![]() |
Open with DEXTER |
![]() |
Figure 3: (continued). |
Open with DEXTER |