A&A 446, 1095-1105 (2006)
DOI: 10.1051/0004-6361:20053951
P. Blay1 - I. Negueruela2 - P. Reig3,4 - M. J. Coe5 - R. H. D. Corbet6,7 - J. Fabregat8 - A. E. Tarasov9
1 - Institut de Ciència dels Materials, Universidad de Valencia, PO Box 22085, 46071 Valencia, Spain
2 - Departamento de Física, Ingeniería de Sistemas y Teoría de la Señal, EPSA, Universidad de Alicante, PO Box 99, 03080 Alicante, Spain
3 - Foundation for Research and Technology-Hellas, 711 10 Heraklion, Crete, Greece
4 - Physics Department, University of Crete, 710 03 Heraklion, Crete, Greece
5 - School of Physics and Astronomy, Southampton University, Southampton SO17 1BJ, UK
6 - X-ray Astrophysics Laboratory, Code 662, NASA/Goddard Space Flight Center, Greenbelt, MD 20771, USA
7 - Universities Space Research Association
8 - Observatori Astronómic, Universidad de Valencia, PO BOX 22085, 46071 Valencia, Spain
9 - Crimean Astrophysical Observatory, Nauchny, Crimea, 334413, Ukraine
Received 29 July 2005 / Accepted 5 October 2005
Abstract
We present the results of our long-term monitoring of
BD
2790, the optical counterpart to the X-ray source 4U 2206+54.
Unlike previous studies that classify the source as a Be/X-ray binary, we find that its
optical and infrared properties differ from those of typical Be stars: the variability of the V/R ratio is not cyclical; there are variations in the shape and strength of the H
emission line on timescales
less than 1 day; and no correlation between the EW and the IR magnitudes or colors is seen. Our observations
suggest that BD
2790 is very likely a peculiar O9.5V star.
In spite of exhaustive searches we cannot find any significant modulation in any emission line parameter
or optical/infrared magnitudes. Spectroscopy of the source extending from the optical to the
K-band confirms the peculiarity of the spectrum: not only are the He lines stronger than expected for
an O9.5V star but also there is no clear pattern of variability. The possibility that
BD
2790
is an early-type analogue to He-strong stars (like
Ori C) is discussed.
Key words: stars: early-type - stars: emission-line, Be - stars: magnetic fields - stars: individual: BD
2790
4U 2206+54, first detected by
the UHURU satellite (Giacconi et al. 1972), is a weak persistent
X-ray source. It has been observed by Ariel V (as 3A 2206+543;
Warwick et al. 1981), HEAO-1
(Steiner et al. 1984), EXOSAT (Saraswat & Apparao
1992), ROSAT (as 1RX J220755+543111; Voges et al. 1999), RossiXTE
(Corbet & Peele 2001; Negueruela & Reig 2001, henceforth NR01) and INTEGRAL
(Blay et al. 2005). The source is variable, by a factor >3 on timescales
of a few minutes and by a factor >10 on longer timescales (Saraswat
& Apparao 1992; Blay et al. 2005), keeping an average luminosity around
for an assumed distance of
(NR01).
The optical counterpart was identified by Steiner et al. (1984), based
on the position from the HEAO-1 Scanning Modulation Collimator, as the early-type
star BD
2790. The star displayed H
line in
emission with two clearly differentiated peaks, separated by about 460 km s-1. Even though some characteristics
of the counterpart suggested a Be star (Steiner et al. 1984), high resolution
spectra show it to be an unusually active O-type star, with an
approximate spectral type O9Vp (NR01).
RossiXTE/ASM observations of 4U 2206+54, show the X-ray flux to be modulated with a period of approximately 9.6 days (see Corbet & Peele 2001; Ribó et al. 2006). The short orbital period, absence of X-ray pulsations and peculiar optical counterpart make 4U 2206+54 a rather unusual High-Mass X-ray Binary (HMXB). The absence of pulsations indicates that the compact companion could be a black hole. Recent studies of high energy emission from the system, however, suggest that the compact object in 4U42206+54 is a neutron star (Blay et al. 2005; Torrejón et al. 2004; Masseti et al. 2004).
In an attempt to improve our knowledge of this system, we have collected optical and infrared observations covering about 14 years.
We present data obtained as a part of a long-term monitoring campaign
consisting of optical and infrared spectra, infrared and optical
broad-band photometry and narrow-band Strömgren optical photometry of
BD
2790, the optical counterpart to 4U 2206+54.
We have monitored the source from 1990 to 1998, using the 2.5-m
Isaac Newton Telescope (INT) and the 1.0-m Jakobus Kapteyn Telescope
(JKT), both located at the Observatorio del Roque de los
Muchachos, La Palma, Spain, and the 1.5-m telescope
at Palomar Mountain (PAL). We have also made use of data from
the La Palma Archive (Zuiderwijk et al. 1994). The archival data
consist of H spectroscopic observations taken with the INT
over the period 1986-1990. The two datasets overlap for a few months
and together they constitute continuous coverage of the source
for thirteen years. The older INT observations had been taken with
the Intermediate Dispersion Spectrograph (IDS) and
either the Image Photon Counting System (IPCS) or a CCD camera. All the INT data after 1991 were obtained with CCD cameras.
The JKT observations were obtained using the St Andrew's
Richardson-Brealey Spectrograph
(RBS) with the R1200Y grating, the red optics and either the EEV7 or
TEK4 CCD cameras, giving a nominal dispersion of
1.2 Å. The
Palomar 1.5-m was operated using the f/8.75 Cassegrain echelle
spectrograph in regular grating mode (dispersion
0.8 Å/pixel).
Further observations were taken with the 2.6-m telescope at the Crimean Astrophysical Observatory (CRAO) in Ukraine.
From 1999, further monitoring has been carried out using the 1.52-m
G. D. Cassini telescope at the Loiano Observatory (BOL),
Italy, equipped with the Bologne Faint Object Spectrograph and Camera
(BFOSC) and the 1.3-m Telescope at the Skinakas Observatory (SKI), in
Crete, Greece. From Loiano, several observations were taken using
grism #8, while higher resolution spectra were taken with grism #9 in
echelle mode (using grism #10 as cross-disperser). Other spectra were
taken with the echelle mode of grism #9 and grism #13 as
cross-disperser, giving coverage of the red/far-red/near-IR region (up
to 9000 Å). At Skinakas, the telescope is an f/7.7 Ritchey-Cretien, which was equipped with a
ISA SITe chip CCD and a 1201 line mm-1 grating, giving a nominal dispersion of 1 Å pixel-1.
Blue-end spectra of the source have also been taken with all the telescopes listed, generally using the same configurations as in the red spectroscopy, but with blue gratings and/or optics when the difference was relevant (for example, from Loiano, grisms #6 and #7 were used for the blue and yellow regions respectively).
All the data have been reduced using the Starlink software package FIGARO (Shortridge et al. 1997) and analysed using DIPSO (Howarth et al. 1997). Table 5 lists a log of the spectroscopic observations.
Near-infrared (I band) spectra of BD
2790 have also
been taken with the JKT, INT and G. D. Cassini telescopes.
K-band spectroscopy of BD
2790 was obtained
on July 7-8, 1994, with the Cooled Grating Spectrometer (CGS4) on UKIRT,
Hawaii. The instrumental configuration consisted of the long focal
station (300 mm) camera and the 75 lines mm-1 grating, which gives a
nominal velocity resolution of 445 km s-1 at 2
m
(
). The data were reduced according
to the procedure outlined by Everall et al. (1993).
We took one set of UBVRI photometry of the source on August 18, 1994, using the 1.0-m Jakobus Kapteyn Telescope (JKT). The observations were made using the TEK#4 CCD Camera and the Harris filter set. The data have been calibrated with observations of photometric standards from Landolt et al. (1992) and the resulting magnitudes are on the Cousins system.
We also obtained several sets of Strömgren uvbyphotometry. The early observations were taken at the 1.5-m
Spanish telescope at the German-Spanish Calar Alto Observatory, Almería,
Spain, using the UBVRI photometer with the
filters, in
single-channel mode, attached to the Cassegrain focus. Three other sets were
obtained with the 1.23-m telescope at Calar Alto, using the TEK#6 CCD
equipment. One further set was taken with the 1.5-m
Spanish telescope equipped with the single-channel multipurpose photoelectric
photometer. Finally, one set was obtained with the 1.3-m Telescope at
Skinakas, equipped with a Tektronik
CCD.
Table 1: Strömgren photometry of the optical counterpart to 4U 2206+54. The last column indicates the telescope used. a stands for the 1.5-m Spanish telescope at Calar Alto. b represents the 1.23-m German telescope. c is the Skinakas 1.3-m telescope.
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Figure 1:
Blue/green spectrum of BD
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All observations are listed in Table 1.
Infrared observations of BD
2790 have been obtained with
the Continuously Variable Filter (CVF) on the 1.5-m. Carlos Sánchez
Telescope (TCS) at the Teide Observatory, Tenerife, Spain and the UKT9
detector at the 3.9-m UK Infrared Telescope (UKIRT) on Hawaii.
All the observations are listed in
Table 6. The errors are much smaller after 1993,
when we started implementing the multi-campaign reduction procedure
described by Manfroid (1993).
There is no evidence for variability in what can be considered with
certainty to be photospheric features (i.e., the Balmer lines from H
and higher and all He I and He II lines in the
blue). However, it must be noted that the EW of H
is
2.2 Å in all our spectra (and this value should
also include the blended O II
4350 Å line), which is too low for any
main sequence or giant star in the OB spectral range (Balona &
Crampton 1974). Average values of EWs for different lines are indicated in
Table 2. The main spectral type discriminant for O-type stars is the ratio
He II 4541Å/He I 4471Å. The quantitative criteria
of Conti & Alschuler (1971), revised by Mathys (1988), indicate that
BD
2790 is an O9.5 V star, close to the limit with O9 V.
Table 2:
Measurement of the EW of strong absorption lines (without
obvious variability and presumably photospheric) in the spectrum of BD
2790.
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Figure 2:
Evolution of the H![]() ![]() |
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Representative shapes
of the H
line in BD
2790 are shown in
Fig. 2.
In all the spectra, two emission components appear clearly
separated by a deep narrow central reversal. The absorption component
normally extends well below the local continuum level - which is
usually referred to as a "shell'' spectrum - but in some spectra,
it does not reach the continuum. The red (R) peak is always stronger
than the blue (V) peak, but the V/R ratio is variable.
The first case of observed strong variability happened during 1986, when the profile was observed to have changed repeatedly over a few months from a shell structure to a double-peaked feature, with the central absorption not reaching the continuum level. The second one took place in 1992, when the strength of the emission peaks decreased considerably to about the continuum level. Finally, during the summer of 2000, we again saw line profiles in which the central absorption hardly reached the continuum level alternating with more pronounced shell-like profiles.
Figure 3 displays a plot of the Full Width at Half Maximum (FWHM),
V/R and peak separation ()
of the H
line against its EW, for all the data from the INT.
H
parameters (EW, FWHM, V/R and
)
were obtained for all the datasets
shown in Table 5. Given the very diverse origins of the spectra
and their very different spectral resolutions, it is difficult to compare them all,
as there are some effects which introduce some artificial scattering in the data. This
is the case of the instrumental broadening affecting the FWHM.
At a first approximation we considered that it was not necessary to account for it.
Taking into account the typical spectral resolutions of our dataset -better than 3 Å in
most cases - and the fact that for the majority of our spectra FWHM > 11 Å (and generally
14 Å), the instrumental broadening, a priori, can be considered negligible.
Dachs et al. (1986) found a correlation between H
parameters (FWHM, peak separation,
EW) in Be stars. We fail to see these correlations when the entire set of
spectra is used but they are present when we restrict the analysis to those
spectra taken with the same instrument, see Fig. 3. There
is, however, a large spread in the case of the V/R ratio. Most of the scatter in FWHM may be related
to the larger uncertainties involved when the emission components are small and the line profile is separated.
Red spectra covering a larger wavelength range (such as that in
Fig. 5) show also the He I 6678 Å line
and sometimes the He I
7065 Å line. Like H
,
the He I
6678 Å line typically
displays a shell profile, but the emission peaks are weaker than those
of H
,
while the central absorption component is normally very
deep. Variability in this line is also more frequent than in H
.
The V peak is generally dominant, but the two peaks can be
of approximately equal intensities and sometimes so weak that they
cannot be distinguished from the continuum. Given the apparent different
behaviour of H
and He I
6678 Å lines, it is
surprising to find that there is some degree of correlation between their
parameters, as can be seen in Fig. 4, where EW of both
lines from INT spectra in which both lines were visible are shown.
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Figure 3:
Parameters of the H![]() |
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The upper Paschen series lines are always seen in
absorption and no variability is obvious (see Fig. 5).
The Paschen lines are much deeper and narrower than those observed
in main-sequence OB stars by Andrillat et al. (1995) and rather resemble
early B-type supergiant stars.
However, it must be noted that some shell stars in the low-resolution catalogue of Andrillat et al. (1988) display I-band spectra that share some characteristics with
that of BD
2790.
K-band spectra are shown in Fig. 6. Unlike the OB components of several Be/X-ray
binaries observed by Everall et al. (1993; see also Everall 1995), BD
2790
shows no emission in He I
2.058
m (though the
higher resolution spectrum suggests a weak shell profile). Br
may have some emission component, but is certainly not in emission.
The situation differs considerably from that seen in the K-band
spectrum of BD
2790 presented by Clark et al. (1999), taken on 1996 October. There Br
displays a clear
shell profile with two emission peaks and He I
2.112
m is in absorption. This shows that the shell-like behaviour and
variability extends into the IR.
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Figure 4:
EW of the He I ![]() ![]() |
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Figure 5:
The spectrum of BD +
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The UBVRI photometric values we obtain
are U=9.49, B=10.16, V=9.89, R=9.88 and I=9.55.
The photometric errors are typically 0.05 mag, derived from the
estimated uncertainties in the zero-point calibration and colour
correction. Table 1 lists our
Strömgren uvby measurements.
V measurements in the literature are scarce and consistent with being constant (see
references in NR01). However, our more accurate set of measurements of the V magnitude
(or Strömgren y) show variability, with a
difference between the most extreme values of
mag
(see Table 1), 0.05 mag being also the standard deviation of
all 7 measurements.
From our UBV photometry, we find that the reddening-free
parameter Q (
Q=-0.72(B-V)+(U-B)) is
.
This value corresponds, according to the
revised Q values for Be and shell stars calculated by Halbedel
(1993), to a B1 star.
We have tried deriving the intrinsic parameters of BD
2790 from our Strömgren photometry by applying the
iterative procedure of Shobbrook (1983) for de-reddening. The values
obtained for the reddening from the different measurements agree quite
well to
(one standard deviation) and the
colour (b-y)0 averages to
.
This value corresponds to a B1V star
according to the calibrations of Perry et al. (1987) and Popper (1980).
Our infrared photometry coverage extends for 13 yr and is
much more comprehensive than our optical photometry. The IR long-term light curve is shown in Fig. 7. Data have
been binned so that every point represents the average of all the
nights in a single run (excluding those with unacceptably large
photometric errors).
As can be seen in Fig. 7, the range of variability is not very
large, with extreme values differing by
in all
three bands. Variability seems to be relatively random, in the sense
that there are no obvious long-term trends. The light curves for the three infrared
magnitudes are rather similar in shape, suggesting that the three
bands do not vary independently.
In spite of this, all colour-magnitude plots are dominated by scatter. Moreover, an analysis of the temporal behaviour shows that there is no obvious pattern in the evolution of the source on the H/(H-K) and K/(H-K) planes, with frequent jumps between very distant points and no tendency to remain in any particular region for any length of time.
The only plot in which a clear correlation stands out is the K/(H-K) diagram (see Fig. 8). In principle, one would be tempted to dismiss this correlation as the simple reflection of stronger variability in K than in H, since (H-K) would necessarily be smaller for larger values of K. However a linear regression of Hagainst K also shows a clear correlation, We find a=0.89, b=0.93 and a correlation coefficient of r2=0.64 for K=aH+b. Suspecting, then, that linear correlation must be present in the H/(H-K) plot as well, we also performed a linear regression. In this case we found a very poor correlation.
Equally disappointing is the search for correlations between the
EW of H
and the (J-K) color. Even though our measurements
of these two quantities are not simultaneous, a look at their
respective evolutions (Fig. 9) shows no clear
correlation.
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Figure 6:
K-band spectra of BD
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All the parameters of the H emission line are clearly variable: EW,
FWHM, V/R ratio and peak separation. In the hope that the variation in
any of these parameters could give us information about the physical
processes causing them, we have searched the dataset for
periodicities. The large variety of resolutions, CCD configurations
and S/N ratios present in the data have hampered our attempts at a
homogeneous and coherent analysis. We have made an effort, within the
possibilities of the dataset, to use the same criteria to measure all
parameters on all spectra. We have used several different algorithms
(CLEAN, Scargle, PDM) in order to detect any obvious periodicities,
but with no success. No sign of any significant periodicity has been
found in any of the trials.
Likewise, we have explored possible periodicities in the infrared
light curves. While the J,
H and K magnitudes seem to vary randomly, we find a striking
apparent modulation of the (J-K) colour. Figure 9
shows an obvious trend in the evolution of (J-K), with a suggestion
that the variability (with an amplitude 0.2 mag) may be
(quasi-)periodic over a very long timescale, on the order of
5 y. Unfortunately, this
timescale is too long compared to our coverage to allow any certainty.
We have also folded the data using the period detected in the analysis of the X-ray light curve of 4U 2206+54 (the presumably orbital 9.56-d period, see Corbet & Peele 2001 and Ribó et al. 2006), without finding any significant periodic modulation.
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Figure 7:
Infrared light curves of BD
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Figure 8: Colour-magnitude plots showing the evolution of the infrared magnitudes. The strong correlation seen in the K/(H-K) plane is not a simple reflection of the fact that a brighter K means a smaller (H-K), as the correlation between H and K is also strong. Regression lines are shown as dashed lines. In the first case the correlation coefficient is r(H-K),K=-0.46 and the correlation is significant in a 98% confidence level. In the latter case the correlatino coefficient is rH,K=0.80 and the correlation is also significant at a 98% confidence level. |
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Considering the possibility that the lack of detectable periodicities
in our dataset was due to the varying resolutions and irregular time
coverage, during July 2000 we carried out a more intensive spectroscopic
monitoring of BD
2790. Observations were made from Skinakas
(Crete) and Loiano (Italy). We collected a set of 2 to 5 spectra per
night during two runs: from 17th to 20th July in Skinakas and from 26th to 31st July in Loiano. The instrumental configurations were identical
to those described in Sect. 2.
We fear that one of our objectives, the study of possible orbital
variations, may have been affected by an observational bias. The
presumed orbital period of the source is 9.56 days, probably too close
to the time lag (10 days) between the first observing night at Skinakas
and the first observing night at Loiano. Therefore we have not been
able to cover the whole orbital period. Indeed, the phases (in the 9.56 d cycle) at which the observations from Skinakas were taken, were
almost coincident with the phases during the first four Loiano
nights. For this reason, our coverage of the orbital period extends to
only 60%, which is insufficient to effectively
detect any sort of modulation of any parameters at the orbital period.
Again, we have measured all parameters of the H line, which
are shown in Fig. 10. Contrary to what we saw when
considering the dataset for the 13 previous years, we find some degree of correlation
between EW, FWHM and
,
while V/R seems to vary
independently. Since this correlation between the different line
parameters seems natural, we attribute the lack of correlations within
the larger dataset to the use of data of very uneven resolution and
quality.
We observe obvious changes in the depth of the central absorption core
in the H line, which is seen sometimes reaching below the
continuum level, while in other occasions is above the
continuum (see Fig. 11). Similar behaviour had already been observed
in 1986 (see
Fig. 2, but no further examples are found in our data
sample). Lines in the blue (3500-5500 Å) are much more stable, as is also the case
when the longer term is considered. In this spectral range, the spectra resemble closely
those obtained at other epochs, with weak emission components
visible in He II
4686 Å and H
.
The reddening to BD
2790 can be estimated in
several different ways. Photometrically, from our value of
,
using the correlation from Shobbrook (1983), we
derive
.
An independent estimation can be made by
using the standard relations between the strength of Diffuse
Interstellar Bands (DIBs) in the spectra and reddening
(Herbig 1975). Using all the spectra obtained from the Cassini
telescope (for consistency), we derive
from the
Å DIB and
from the
Å DIB. All these values are consistent with each other,
therefore we take the photometric
value as representative of the reddening to BD
2790.
From five measurements available in the
literature (including the one presented in this work), we find
.
With the E(B-V) derived, this indicates an
intrinsic colour
,
typical of an early-type
star, confirming the validity of the reddening determination. As
discussed in NR01, the value of the absorption column derived from all X-ray
observations is one order of magnitude larger than what is expected from the
interstellar reddening. This affirmation stands also when we consider the more
accurate measurement of the absorption column
(i.e.,
)
from BeppoSax
data (Torrejón et al. 2004; Masseti et al. 2004).
Averaging our 7 measurements of y with the 5 V measurements, we
find a mean value for BD
2790 of
.
Assuming a standard reddening law (R=3.1), we find
V0=8.21. If the star has the typical luminosity of an O9.5V star
(
MV=-3.9, see Martins et al. 2005), then the distance to BD
2790 is
kpc. This is closer than previous estimates (cf. NR01), because
the absolute magnitudes of O-type stars have been lowered down in the most recent
calibrations.
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Figure 9:
Evolution of the infrared colours in BD
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Figure 10:
H![]() |
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Figure 11:
Evolution of H![]() ![]() |
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Since its identification with 4U 2206+54, BD
2790 has
always been considered a classical Be star, because of the presence of
shell-like emission lines in the red part of its spectrum. However,
the main observational
characteristics of BD
2790 differ considerably from
those of a classical Be star:
We estimate that the most likely spectral classification of BD
2790
is O9.5Vp. However some remarkable peculiarities have been noticed:
while the blue spectrum of BD
2790 suggests an 09.5 spectral
type, there are a few metallic lines reminiscent of a later-type spectrum (see NR01); the UV lines
support the main sequence luminosity classification, but the Paschen lines resemble
those of a supergiant.
In order to obtain a measure of the rotational velocity of BD
2790 we
have created a grid of artificially rotationally broadened spectra from that of the
standard O9V star 10 Lac. We have chosen 10 Lac because of its very low projected rotational
velocity and because the spectrum of BD
2790 is close to that of
a O9V star.
In Fig. 12 normalised profiles of a set of
selected helium lines (namely, He I
4026,
4144,
4388, and
4471 Å) are shown together with the artificially
broadened profile of 10 Lac, at 200 km s-1and those rotational velocities producing upper and lower envelopes to the
widths of the observed profiles of BD
2790. The rotational
velocity of BD
2790 must be above 200 km s-1. For
each line, the average of the rotational velocities yielding the upper
and lower envelopes were taken as a representative measurement of the rotational
velocity derived from that line. The results of these measurements are summarised
in Table 4. We estimated the averaged rotational velocity
of BD
2790 to be
km s-1.
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Figure 12:
Normalised profiles of selected He I lines (namely, He I ![]() ![]() ![]() ![]() ![]() ![]() |
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Comparison of the helium profiles with those rotationally broadened from 10 Lac
shows that the observed helium profiles in BD
2790 are stronger than
what is expected for a normal O9.5V star.
The strength of the He lines suggests the possibility that BD
2790 may be
related to the He-strong stars. These are a small group
of stars, with spectral types clustering around B2 V, that show anomalously strong helium lines.
A well known O-type star believed to be related to He-strong stars is
Ori C, which is known to vary in spectral type from O6 to O7
(Donati et al. 2002; Smith & Fullerton 2005). BD
2790
could be the second representative of this class of objects among O-type stars.
He-strong stars display a remarkable set of peculiarities:
oblique dipolar magnetic fields, magnetically controlled winds, and chemical surface anomalies,
among others. Usually these stars are distributed along the ZAMS
(Pedersen & Thomsen 1997; Walborn 1982; Bohlender et al. 1987; Smith & Groote 2001).
A rich variety of phenomena have been observed in these objects: in the UV, they can show
red shifted emission of the C IV and Si IV resonance lines (sometimes variable);
in the optical bands they are characterized by periodically modulated H emission,
high level Balmer lines appearing at certain rotational phases and periodically modulated
variability in He lines, sometimes showing emission at He II
4686 Å. They
can also show photometric variability with eclipse-like light curves.
Table 4: Summary of the measured rotational velocities for the selected helium lines shown in Fig. 12.
Except for the periodic modulation of the variations, BD
2790 shares
many of these peculiarities. In particular, together with the apparent high helium
abundance, BD
2790 shows variable H
emission and
He II
4686 Å emission, the UV spectrum shows apparently prominent
P-Cygni profiles at C IV and Si IV resonance lines (see NR01).
In contrast a wind slower than expected is found (see Ribó et al. 2006), which can be an indication
of some red-shifted excess of emission in these lines. In He-strong stars the wind is
conducted along the magnetic field lines into a torus-like envelope located at the magnetic
equator. This configuration can lead to the presence of double emission peaks in H
,
which
resemble those seen in BD
2790, but which usually show a modulation on
the rotational period.
The complexity and shape of the double peak will depend on the angle between magnetic
and rotational axes and the line of sight to the observer (see Townsend et al. 2005).
A rotationally dominated circumstellar envelope is clearly present in BD
2790, as indicated by the infrared magnitudes,
the emission in Balmer and some helium lines and the correlations between H
line parameters. However the structure of this circumstelar envelope clearly differs from those seen in Be stars.
Following the analogy with He-strong stars, the existence of a circumstellar disk-like structure is also common to
this type of objects. The only difficulty to accept BD+53
2790 as a He-strong star
is the apparent lack of rotational modulation of the emission lines parameters. Given the rotational
velocities derived, we could expect a rotational period of a few days. In addition to the problems
in the diverse origin of our data (see Sect. 3.1), the sampling on time of our measurements is not adequate to find
variations on time scales of a few days (modulated with the rotational period), thus we cannot discard
yet the presence of rotational periodicity. The idea of a magnetically driven wind contributing to a
dense disk-like structure is not strange even in the modelling of Be stars' circumstellar envelopes. The wind
compressed disk of Bjorkman & Cassinelli (1992) was shown to be compatible with observations
only if a magnetic field on the order of tens of Gauss was driving the wind from the polar
caps onto the equatorial zone (Porter 1997).
A careful inspection to the correlation seen in Fig. 4 between the
He I 6678 and H
EWs shows that there is a common
component to the emission of both lines. H
emission, then, will have at least two contributions:
a P-Cygni like contribution (as seen in the 1992 spectra, see Fig. 2, where the double peak
structure disappears and only the red peak survives) and an additional variable double
peaked structure. The relative variation of both components may hide any
periodic modulation present.
Therefore, we can conclude that this is a very peculiar O9.5V star where most likely a global strong magnetic field may be responsible for most of the behaviour seen so far.
We have presented the results of 14 years of spectroscopic
and optical/infrared photometric monitoring of BD
2790, the optical component of the Be/X-ray binary 4U 2206+54.
The absence of any obvious long-term trends in the evolution of different parameters and
fundamentally the absence of correlation between the EW of H
and the infrared
magnitudes and associated colours makes untenable a Be classification for the star. Based on a careful inspection
to the source spectrum in the classification region and the peculiar behavior of the H
emisson line, we conclude
that the object is likely to be a single peculiar O-type star (O9.5Vp) and an early-type analogue
to He-strong stars.
Acknowledgements
We would like to thank the UK PATT and the Spanish CAT panel for supporting our long-term monitoring campaign. We are grateful to the INT service programme for additional optical observations. The 1.5-m TCS is operated by the Instituto de Astrofísica de Canarias at the Teide Observatory, Tenerife. The JKT and INT are operated on the island of La Palma by the Royal Greenwich Observatory in the Spanish Observatorio del Roque de Los Muchachos of the Instituto de Astrofísica de Canarias. The 1.5-m telescope at Mount Palomar is jointly owned by the California Institute of Technology and the Carnegie Institute of Washington. The G. D. Cassini telescope is operated at the Loiano Observatory by the Osservatorio Astronomico di Bologna. Skinakas Observatory is a collaborative project of the University of Crete, the Foundation for Research and Technology-Hellas and the Max-Planck-Institut für Extraterrestrische Physik.
This research has made use of the Simbad database, operated at CDS, Strasbourg (France), and of the La Palma Data Archive. Special thanks to Dr. Eduard Zuiderwijk for his help with the archival data.
We are very grateful to the many astronomers who have taken part in observations for this campaign. In particular, Chris Everall obtained and reduced the K-band spectroscopy and Miguel Ángel Alcaide reduced most of the H
spectra.
P.B. acknowledges support by the Spanish Ministerio de Educación y Ciencia through grant ESP-2002-04124-C03-02. I.N. is a researcher of the programme Ramón y Cajal, funded by the Spanish Ministerio de Educación y Ciencia and the University of Alicante, with partial support from the Generalitat Valenciana and the European Regional Development Fund (ERDF/FEDER). This research is partially supported by the Spanish MEC through grants AYA2002-00814 and ESP-2002-04124-C03-03.
Table 3: Log of spectroscopic observations during 2000.
Table 5: Log of spectroscopic observations. Some representative spectra are displayed in Fig. 2 (marked with *).
Table 6: Observational details, IR photometry.