Astron. Astrophys. 333, 125-140 (1998)
4. Phases of a line emission outburst
4.1. Relative quiescence
According to Fig. 3, the emission strength of µ Cen
is changing most of the time. Therefore, one can only define phases of
relative quiescence when the variations are minimal. Even during such
relative quiescence the Balmer emission slowly decays. This definition
is independent of the emission strength.
In our observations, we have found the end of such a phase in 1996
at Modified Julian Date (MJD JD-2 400 000.5)
, nearly 80 days after the last major event had
occured. We will base the following description on this period, but
the spectra during the relative quiescence before MJD 49 785 look
essentially the same. The small differences might be explained by the
shorter time that had elapsed since the preceding outburst.
- Balmer lines: At this stage, the emission profiles have
sharply defined edges, and the central reversal is very pronounced.
The -ratio is only slightly variable about
The contribution of extended emission to the line wings is quite weak,
so that the photospheric absorption wings are still visible. The
process responsible for the extended wings might be Thomson scattering
in this phase. The quiescent -line is shown in
- Paschen lines: The emission in the Paschen series is
shallow in the spectral range observed (Paschen discontinuity to
Pa14). The peak height of Pa15, the blue side of
which is slightly blended with Ca ii 8542, does
not exceed = 1.1. This level of the Paschen
emission during quiescent phases may be crudely recurrent whereas the
strength of the Balmer emission during quiescence takes on a much
larger range of values. The central reversal is much less pronounced
than in the Balmer series, though detectable. Apart from the Paschen
lines, three other lines within the Bracket continuum were detected in
the observed spectral range: Ca ii 8498,8542
and O i 8446. The first two are weak and
moreover blended with Paschen lines. But the latter is stronger than
one would expect judging from the O i 7774
blend. It is a fluorescence line of the
transition (Briot, 1981), and its strength should, therefore, scale
with . The quiescent Pa15 line
profile is shown in Fig. 5.
- Silicon: Weak Si ii 6347 emission is
present (Fig. 5) whereas Si ii 4131,5056
are missing. Of the two main phases of relative quiescence in 1995 and
1996, the later one was weaker in Si ii 6347
emission. As can be seen in Fig. 6, also the peak separation was
- Iron: During relative quiescence, Fe ii emission is not
detectable down to the noise limit which, given the large number of
potential lines, corresponds to a peak height of half a percent above
the continuum level.
- Helium: The emission profiles of He i are highly variable
in both shape and strength even at quiescence. However, the absolute
level of activity and the velocities of the emission peaks are
generally lower than during an outburst, and the emission peaks are
narrower than in other stages.
4.2. Precursor phase
An imminent outburst announces itself by a short (5-13 days)
significant (see Sect. 3and also Fig. 1) drop in the peak
height of all emission lines present at the time. A second defining
constituent of this phase is an increase of the extended emission
wings which is best visible in higher Balmer lines. In
these wings may compensate the change in
equivalent width caused by the drop of the peak height, however not
quite in all cases. The balancing is more perfect in H
, and from H on the wings
grow more strongly than the peaks decrease.
In most outburst cycles, this phase is the most distinctly
recognizable feature; a proto-typical example is displayed in
Fig. 3 around MJD 50 510, but the reduced
-peak height is also well visible in Fig. 5. This behaviour is
not peculiar to µ Cen as is shown by the case of the
Be star HD 76 534 in which Oudmaijer & Drew (1997) observed
the full recovery of a previously drastically reduced H
emission in only 3 hours.
The variations of (leftmost), Pa15 (mid left), Si ii 6347 (mid right), and He i 6678 (rightmost), from relative quiescence (lowermost spectrum, MJD 50 180) over the very first trace of circumstellar variability, the beginning of the precursor (second, MJD 50 184), the high velocity absorption event as described in Sect. 6(third, MJD 50 185), and the appearance at the end of the burst phase (, fourth, MJD 50 195) to the late relaxation phase (, fifth, MJD 50 204). The dates of the spectra are indicated as arrows in Fig. 3. Two similar high-velocity absorption events (Table 4 and Sect. 6) in January 1997 are overplotted as dotted lines on He i 6678 (MJD 50 457 in the lower and MJD 50 461 in the upper spectrum). In the lower row we show how the dynamical spectra develop from relative quiescence through a major burst into the relaxation phase. The dynamical spectra also cover two smaller bursts around MJD 50 223 and MJD 50 232. The apparent long-term radial-velocity variations in the He i 6678 line are due to the highly uniform sampling of one spectrum per day, which results in a strong beat pattern with the photospheric periodicity we describe in Paper II. The feature in the Pa15 line seen in the dynamical spectrum at is a CCD artifact
The peak separation of the Si ii 6347 () and Fe ii () lines. For 1995 (top), the peak separations of Fe ii 5169 and Fe ii 5317 are averaged. For 1996 (bottom), only Fe ii 5169 was measurable with sufficient accuracy. Note that Fe ii emission becomes detectable only at the peak of an outburst (Fig. 2)
Because no photometry was obtained parallel to the
HEROS observations, the possibility cannot be directly
dismissed that the precursor drop of the ratio
is caused predominantly by a corresponding increase of the continuum
flux. In the precursor phase of the outburst on MJD 50 515, which was
the deepest observed, the H
value temporarily dropped from 3.0 to 2.4. If
the decrease was caused only by a change of the continuum flux, a
brightening by nearly 0. would be implied. This
is more than virtually all optical peak-to-peak light variations
reported by Cuypers et al. (1989), Balona (private communication), and
the HIPPARCOS photometry (Perryman et al., 1997) for the period 1987
to 1992. Taking into account also the frequency of the outbursts and
the density of the photometric data, such a large change in the
continuum flux appears improbable.
Another test is to check the amplitudes of
various lines for consistency with pure continuum variability.
Unfortunately, there are too few metal lines that are of sufficient
strength throughout the outburst cycle and to which the test can be
applied. However, accompanying major Balmer decrement variations
(Rivinius et al., in preparation) indicate that the results of this
test, too, would be negative.
Baade et al. (1988) present arguments that previous rapid increases in
line emission from µ Cen were due to the ejection of
material by the star to its circumstellar disk. This justifies the
notion of outbursts. Hanuschik et al. (1993) distinguish between
major and minor outbursts which can be identified by their strength.
We find two major and many minor bursts in our data but no compelling
evidence that they are genuinely different events. The difference
between bursts of different strengths may be blurred further by the
dependency of the appearance of bursts on the mean emission strength.
Outbursts tend to look more uniform at times of stronger underlying
disk emission, and we note that Hanuschik et al. derived their
classification from data obtained when there was not yet a new
Of the two major outbursts, the one in 1995, the major burst
unfortunately started shortly before a gap in our observing schedule
(cf. the description of the behaviour of the Fe ii lines given below).
Both major outbursts were sampled with one spectrum per day. At higher
temporal resolution only minor bursts were observed with
HEROS ; the series of Boller&Chivens spectra was
probably also obtained during a minor outburst (Sect. 5.2).
- Balmer lines: In 1996, the peak height of
dropped after the short precursor phase within
a few days by 0.3 in units of the local continuum, while the
equivalent width change is roughly compensated by the increasing wings
(indicated by the dotted line in Fig. 3, 1996 panel). A
comparison between different Balmer lines shows that the increase of
the wings is not correlated with the strength of the central emission.
Therefore, Thomson scattering is questionable as a major contributor
to the emission wings. The rapid variability may
occasionally have started already even in the precursor phase,
depending on the strength of the burst. The emission wings still
continued to rise and, depending on the level of the quiescent
emission, may finally also in overcompensate the
decrease in emission strength caused by the loss in peak height.
During the 1995 HEROS observations a major burst
occurred, but unfortunately very close to a gap in the observing run,
when the instrument was for two weeks off the telescope. Nevertheless,
the onset of the burst on MJD 49 801 can be detected both in the
equivalent widths (the dotted line in Fig. 3, 1995 panel) and in
the dynamical spectra.
Owing to the increasing optical thickness of
during the past few years, the distinction between major and minor
bursts becomes on the basis of the variation of the H
emission less clear than before when only a
non-persistent emission disk surrounded the star. However, this
classification can still be retained if applied to other lines within
the HEROS wavelength range, for instance to those of Fe
ii described below.
- Paschen lines: The Paschen series exhibits the most drastic
changes during an outburst. The emission strengthens by several
Å, which is due not only to a strengthening in peak height, but
also to strongly enhanced wings. The central reversal has nearly
disappeared and the peak height of the line is of the order of
. The O i 8446 line is not
showing this strong variability. It has nearly not changed its peak
height and is only slightly broader than before so that after an
outburst it is weaker than the neighbouring Pa18 line,
whereas it is stronger during quiescence. Similarly to
, major and minor outbursts are distinguishable
by the amplitude of the increase in emission as long as the persistent
emission is not too strong. This was the case in 1995 and 1996, while
in 1997 the distinction was less clear.
- Silicon: The Si ii-emission becomes visible in
4131,5056 and strengthens for
6347. The peaks are at higher velocities
compared to the stage of relative quiescence. In Si ii
6347 also remnants of the emission from the
previous outburst are still visible at lower velocities. This
behaviour is displayed in Fig. 5 as a dynamical spectrum and is
reflected in the peak separation which seems to reach high values
immediately (Fig. 6). It can also be seen that major outbursts
are far more efficient in driving up the peak separation than are
- Iron: At the peak of major outbursts, the Fe ii emission
appears rather suddenly and within 10 days attains its maximum
strength (cf. Fig. 2). It then persists throughout the outburst
phase and shows a similar behaviour as the SiII
-emission. The distinction between major and minor outbursts is
clearest for these iron lines: only a major burst leads to persistent
emission, while the emission of a minor burst, if present at all,
vanishes as the burst ceases.
- Helium: In He i 6678 additional
emission shows up at relatively high velocities. The differences
between major and minor bursts are less strong in helium than they are
in lines formed farther away from the star.
Finally, it should be pointed out that the pre-outburst loss in
emission peak height (Sect. 4.2) is not always quite compensated
for during the outburst and that the recovery also takes part of the
relaxation phase (Sect. 4.4). In rare cases, a small net loss
may remain well into the next outburst cycle.
In the two small bursts observed with HEROS at
sufficiently high temporal resolution (see Sect. 5) we find that
the variability is most pronounced on the
ascending branch of the emission strength curve. We therefore adopt
the settling of the rapid -variability as the
end of a burst and the beginning of the following phase to which
hereafter we refer as the relaxation phase.
- Balmer lines: The wings of the emission lines attain
their maximum strength at the beginning of this phase. No trace of the
photospheric profile is detectable.
The decline of the wings sets in slowly while the emission peak height
finally increases. This increase may continue for weeks, eventually
turning over into a slow decrease when relative quiescence
Although the -variability has stopped, the ratio
is still not unity. There is a general trend for the red peak to
remain lower than the blue one by several percent. Only after 10 to 20
days (for major outbursts), the average -ratio
finally reaches unity. Line profiles that are representative of these
stages are provided in Fig. 5.
- Paschen lines: They are initially also still dominated by
the broad wings (Fig. 5). During the course of a couple of days
the wings fade slowly, the peak height decreases quite linearly, and
so the shapes of the profiles asymptotically approach their quiescence
- Silicon and iron: Fig. 6 shows the measured peak
separation for Si ii 6347 and two Fe ii lines.
The separation attains high values immediately after a major burst and
then shrinks until the next major burst. Subsequent minor bursts only
have little influence on the peak separation. An example of this
behaviour is given in Fig. 2 in the form of a dynamical spectrum
for Fe ii 5169.
- Helium: Except for a further weakening and narrowing of the
emission components and ongoing variability, the
helium lines seem to be the first to reach the quiescent phase.
© European Southern Observatory (ESO) 1998
Online publication: April 15, 1998