As we have seen dramatic changes in the photometry of HD 45677 on different
time scales one could expect to see such changes in the spectroscopic
behaviour as well. Indeed, spectral variations have been observed for the
lines of MgII, HeI and slightly in CaII
(Swings et al. 1980).
Because of the high radial velocities measured for the
HeI and MgII lines, 150-200 km s
, those lines are
thought to be formed close to the stellar surface. For the H
and
H
lines more dramatic variations are seen as well, most significantly
in the violet parts of the profiles and on time scales ranging from hours up
to days (Merrill 1952;
Swings et al. 1980). Similar variations in the
H
profiles were already reported on a longer time scale on plates
taken between 1923 and 1947 (Swings & Struve 1943;
Merrill 1952). Although
Swings & Allen (1971)
did not report any correlations of line variations
with detected photometric changes, we can give a more exact view at the
moment since more spectroscopic observations overlap with the different time
scales of the photometric variations.
To study the long time scale variations we use as a starting point the
spectral descriptions of Merrill (1928)
and of Swings et al. (1980).
We have searched for additional spectra which are listed in Table 1 (click here).
This summary is certainly not complete as probably many plate and CCD spectra
are reported but not discussed, e.g.
Jaschek et al. (1992), or unpublished,
e.g. Swings (1995). Several of such spectra, as well as many others listed
in Table 1 (click here), will be discussed in detail by
Israelian et al. (1996).
We refer to the spectra in the following sections by their
number, the year of observation, as listed in Col. 1 of Table 1 (click here).
Most historical known spectra are of low resolution. It is therefore only
possible to make use of our own low and intermediate resolution spectra for
these purposes. As our low resolution spectra are quite similar we have only
presented the 1992e and 1992g in Fig. 4 (click here), together with the line identifications
which will be discussed. For a similar reason the 1992b intermediate resolution
spectrum is not shown in Fig. 5 (click here). However, from the presented H
profiles
clear changes are seen.
Investigating such short time scale variations, as also known from the
literature, we use our monitoring data. We have obtained time series with a
duration of 6 days by high resolution spectroscopy in two wavelength regions.
Besides spectra centered on the H
laboratory wavelength, with a
resolving power of 50 000, we have monitored the NaID lines, with
a resolving power of 55 000. The choice for this line is inspired by the
clear detection of accreting cool material in the case of the Herbig Ae star
UX Ori (Grinin et al. 1994). Furthermore, our setup includes the HeI
line at 5876 Å and the [FeII] line at 5870 Å.
The NaID and H
time series were taken in October 1993.
One other spectrum, in both regions, was obtained in January 1995. The
spectra are plotted in Figs. 9a and 10a.
There is no spectral data that would illustrate changes on intermediate time scales in either low or high resolution. This will therefore not be discussed.
The existence of an extended gaseous shell of considerable density was first inferred from the appearance of nebular lines of [OI] at 6300 Å and 6363 Å and of [SII] at 4068 Å (Merrill & Burwell 1933; Swings & Allen 1971).
Contrary to spectroscopic variations as mentioned above, the forbidden lines, like [FeII], [OI], [NII], are stable; no clear changes are seen since the first detections. Probably those lines are formed in an extended region far away from the central star, see Fig. 13 (click here) of Swings (1973). The red [SII] lines were not detected and the [NII] lines very weak. As reported by Swings (1973), the FeII lines are double peaked and more dispersed than the [FeII] lines.
Because of the wide variety of the spectroscopic data in wavelength coverage and resolution we will discuss them by each group separately. In some cases a more detailed description will be given by Israelian et al. (1996).
Figure 4: December 1992 low-resolution spectra of HD 45677 with the most
prominent lines identified
Table: Identified emission lines in the low- and intermediate-resolution
spectra of HD 45677, Figs. 4 (click here) and 5 (click here).
Uncertain identifications are given between brackets
In the 1971 and 1992 spectra, H
is seen in emission accompanied at
the red side by the [NII] line. In the intermediate resolution
1992a and b spectra H
is also accompanied at the blue side by the
weaker [NII] line of the same multiplet. Both are in emission.
It is also seen in most other spectra, like
in 1989, in which H
shows self-absorption. The velocities and line
profiles measured in the 1989 spectrum are similar to those reported before,
e.g. Merrill (1952). However, the H
profile is very complex. We
will discuss this further in Sect. 4.2.2.
In the 1992c-e spectra, as in the previous ones, the OI lines at
7774 Å and 8446 Å are very strong. The [SII] lines, which
were not detected before, are weakly in emission (spectra 1992a and b).
No differences can be seen for the FeII lines in the 1992 spectra
and those observed before. The H
and other blue hydrogen
line profiles confirm the detections given by
Swings (1973) and
Swings et al. (1980).
From 1964 to 1966 emission components appeared in the blue
hydrogen lines, from H10 up to H5 (= H
). In 1979 H
shows
central emission. In 1974, Swings (1974) reported that H
exhibits a
very strong P Cygni profile. In 1966 no emission is seen.
Furthermore, the known strong iron emission lines were observed
double peaked, separated by 
30 km s
,
being the same value as reported above.
In the 1992a and b spectra H
is seen in emission with a central
emission embedded in a photospheric absorption line, as the profile was
observed in the 1971 spectra. In the 1992 intermediate resolution spectra
we also note that the FeII lines are relatively broad, when compared
to typical nebular lines as [NII], [SII] and [OI]
in the same spectra. The FeII lines also hide traces of double peaked
profiles with a separation of about 40-60 km s
, which
is somewhat higher than in the spectra taken before the photometric
minimum. Although the FeII lines are relatively broad and the
resolution of the 1992 spectra are too low to determine good
velocities, most identified lines seem to be centered on a velocity of
about 20 km s
.
Figure 5: May 1992 intermediate-resolution spectra of HD 45677. The full
wavelength range of the 1992a spectrum is shown, the 1992b spectrum is very
similar and therefore not included. However, line profiles can be very
different as shown in the lower plots for H
The properties of the described spectra are similar to those observed and
discussed by Swings (1973) and
Swings et al. (1980). Even the short time
scale variations reported in these papers are seen in the 1992 intermediate
resolution spectra. We have compared both of the spectra, 1992a and 1992b.
The difference spectrum (not shown here) clearly indicates significant
changes in the H
line while all other lines seem to be stable. In
the 1992b spectrum we clearly have a broad extra absorption at the blue side
of the H
spectrum, as well as two small absorption components at +4
and -60 km s
. In the next section we will discuss the H
profile in more detail.
High resolution spectra of HD 45677 are rare. A comparison of our data with
other spectra, besides the hydrogen lines, is therefore not possible.
However, we will discuss here the spectra mentioned in Table 1 (click here) of
CaIIK (Fig. 6 (click here)) and [OI]
6300 (Fig. 7 (click here))
including the MgI 6318.23/.75 Å doublet (Fig. 8 (click here)). The other high
resolution spectra will be discussed in the next section.
A H
spectrum from Doazan et al. (1991) shows a strong double peaked
emission line with the red component much stronger than the violet peak
similar to the 1989a spectrum and the ones reported before
(e.g. Merrill 1928). However, the peak separation in the H
spectra by
Doazan et al. (1991) and by
Merrill (1928) is much bigger (150 and 170
km s
, respectively) than in all the H
spectra (
km s
). This difference supports our interpretation of the double-peaked
structure of the hydrogen lines as being due to self-absorption rather than
being due to two separate regions of hydrogen emission. For our 1993 H
spectra the two peaks differ less in strength than in the 1988b H
spectrum, while in the 1992a and b H
spectra we even have 
. This indicates a strong variability of the hydrogen lines as
noted previously by Swings et al. (1980).
The CaIIK line is seen with the absorption center at
20 km s
. At the blue and red side of this narrow line (the
maximum velocity dispersion, at the continuum level,
is less than 50 km s
) weak emission is seen. This profile
is identical to the ones from 1932 to 1947 (Merrill 1952)
and similar to the
one mentioned by Swings (1973). The
[OI]
6300 and the MgI doublet are in emission with
the central peak at about 20 km s
. Although the velocity dispersion of
the [OI] line is higher, about 90 km s
, than the CaII line,
it is much lower than the 250 km s
of the MgI doublet. Note
also that the latter is asymmetric with extra emission at the blue side of its
profile extending to 150 km s
, whereas in the red part the velocity
goes up to about 100 km s
.
From the velocities we conclude that the MgI line is formed
closer to the central star than the [OI]
6300, which must
be formed in an extended low density region. Although MgI will be
destroyed close to the central star there must be some outflow of this
material.
Also the emission in the CaII resonance line shows the presence of
some cool material. When distributed in a disk this material must be very
dense as seen by the strong self-absorption line.
In comparing the red spectra of 1971b and 1973, although the 1971b spectrum is of poor quality, we see in the 1971b spectrum HeI 6678 Å in emission but not in the 1973 spectrum. In 1973 FeII 7712 Å and OI 7774 Å are in emission but in 1971b weak or probably not. In the red part of 1992c and in 1992d the FeII 7712 Å, [FeII] 7732 Å and OI 7774 Å are seen in emission, similar to the 1973 situation. In 1993 these lines are also visible in emission, only the OI line was weaker than it was in 1992d. In the 1989b spectrum the OI 8446 Å is strong in emission, also the Paschen lines and many FeII lines are in emission, very similar to the 1973 situation. The same wavelength region of the 1992d spectrum looks exactly like in the previous ones, with the exception of the one detecting the decreasing strength of OI 7774 Å.
The 1988 spectra look exactly as the ones of HD 87643 and CD-4211721, two other well known B[e]-stars. Only H-lines are in emission, no other features are seen. No other corresponding spectra are available for comparison.
The high resolution sodium and H
line profiles are presented in Fig. 9 (click here)a
and
Fig. 10 (click here)a, respectively. Their characteristics will be presented separately.
A more qualitatively modeling on these spectra will be presented in
another paper (Israelian et al. 1996).
It is striking to see that in all spectra the NaID lines are
significantly in emission with a strong absorption core at 20 km s
,
similar to the observed CaIIK line.
The emission components of the 
lines is somewhat stronger than the
line, as usual, but combined with the relatively strong
self-absorption component the 
profile looks somewhat different.
The interstellar component is not visible and is probably blended in the
broad circumstellar absorption feature.
Between the spectra taken in 1993 and the one of 1995, no dramatic changes
are seen, although the strong telluric waterlines could give a wrong
impression. However, comparing the time series spectra to each other,
using the difference between a spectrum and the one of 09/10/1993,
it is clear that day by day variations are significant in both the
NaID and the HeI lines (Fig. 9 (click here)b). Note the changes of
the emission strength at the
blue sides of the NaID lines. However, the wings of the
red sides of these lines are remarkably stable as they all
extent up to 147 
5 km s
for the 
and up to
140 
5 km s
for the 
components.
Similar profiles, but with more dramatic changes, are observed by
Israelian et al. (1996) on a somewhat longer time scale.
Figure 6: High resolution CaIIK spectrum of HD 45677, rebinned to
heliocentric velocities
Figure 7: High resolution [OI]
6300 spectrum of HD 45677,
rebinned to heliocentric velocities
Figure 8: High resolution spectrum of the MgI 6318.23 + 6318.75 Å
doublet of HD 45677, rebinned to heliocentric velocities
In the spectra of 13 and 14/10/1993
the blue emission wings are quite strong and extend up to
-99 
km s
and -65 
5 km s
, respectively, while on 10
and 11/10/1993 these blue parts are depressed, and show even weak narrow
absorption components at -18 
1 km s
(not to be
confused with the telluric waterline as seen in the 16/1/1995
spectrum). On 09 and 12/10/1993 these weak lines are located at
-11 
1 km s
. An extra wide but very weak
absorption is also seen on 11/10/1993 at -75 
2 km s
.
Figure 9: a) High resolution NaID spectra of HD 45677, rebinned
to heliocentric velocities. b) The same spectra but taken as the difference
compared to the one of 09/10/1993
The blueshifted weak absorption components seem to be dynamical on a daily time scale, but can also appear for a few days, with relatively large changes in their velocities. In our time series this absorption slowly disappears and in the last two spectra, 13 and 14/10/1993, only extra emission is seen.
lines
Several H
spectra were taken in the same nights as the
NaID spectra. The one taken on 11/01/1994 is somewhat saturated,
as is the 09/10/1993 spectrum slightly overexposed.
As we noticed only small daily changes between the NaID spectra,
the same seems to be the case for the quasi-simultaneously taken H
profiles, see Fig. 10 (click here)a. The H
profiles are quite complex.
A strong self-absorption component
at +11 
1 km s
is visible in all the spectra. However, it is
not located at the center of the emission profile.
We have determined the center of the emission profile by the fitting of a
Gaussian to the observed profiles and to modified ones, in which the
absorption components are deleted by interpolation. The results of both
fitting procedures are almost identical. The emission peak seems to shift
from +37.5 
1 km s
for the 09/10/1993 profile to
+40.5 
1 km s
for the 12/10/1993 spectrum, with a daily step
of +1 km s
. This is almost 10 km s
higher than the radial
velocity of the H
lines before 1928
(Merrill 1928).
The velocity offset between the
H
emission and systemic velocity from +17.5 to +20.5 km s
indicates gaseous movements and acceleration towards the central star.
It would be interesting to monitor HD 45677 daily
for a longer time-interval to reconfirm such a mechanism.
Figure 10: a) High resolution H
spectra of HD 45677, rebinned to
heliocentric velocities. b) The same spectra but taken as the difference
compared to the one of 09/10/1993. Because of saturation effects the spectrum of
11/01/1994 is not shown here
In a plot of the differences between the spectra (Fig. 10 (click here)b), the daily movement
is seen as the sharp extra emission on the red side. Furthermore it is clear
that an extremely broad and very weak component in the H
profile
changes from a slightly redshifted absorption (10/10/1993), in fact less
emission as it is the difference spectrum with the one taken on 09/10/1993,
to extra blueshifted emission (11 and 12/10/1993), note the velocity
dispersion close to the continuum level in these spectra.
Again, this daily change visible as small V/R-variations,
with an extremely broad velocity dispersion, is
either due to the underlying photospheric line or to variations in the wings
of the emission profile. Anyway, it seems that infalling gaseous material is
followed by extra outflow. This process must be studied further, but it will
be related to a region close to the stellar surface.
Another important detection in the H
profiles are the
weaker absorption components and very weak dips, besides the very strong
absorption line at +11 km s
.
We detected these components at -80, -37, +49 and +90 km s
.
For the weaker ones the error is about 5 km s
.
The redshifted components are more prominent.
Table 4: SED fit results for HD 45677 at different V magnitudes
Figure 11: Typical observed (squares) and extinction-corrected (circles)
SEDs for HD 45677. Also shown is the Kurucz model for a B2 V star, fitted
to the extinction-free SED
