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\IBVShead{5013}{15 January 2001}

\IBVStitle{H\boldmath{$\alpha$} Observations of T~C\lowercase{r}B}

\IBVSauth{Zamanov, R.$^1$; Marti, J.$^2$}

\IBVSinst{National Astronomical Observatory ``Rozhen'', P.O. Box 136, BG-4700 Smolyan, Bulgaria\\ \indent e-mail: rozhen@mbox.digsys.bg}
\IBVSinst{Departamento de F{\'{\i}}sica, Universidad de Ja{\'e}n, Spain,
C/ Virgen de la Cabeza 2, 23071 Ja{\'e}n, Spain\\ \indent e-mail: jmarti@ujaen.es}

\SIMBADobj{T CrB}
\IBVStyp{ symbiotic }
\IBVSkey{photometry}
\IBVSabs{Halpha spectroscopic observations of the recurrent nova T CrB during}
\IBVSabs{1993-2000 are reported. A new maximum of the EW(Halpha)}
\IBVSabs{is observed between JD 2450600-JD 2450900.}
\IBVSabs{Very weak traces of the orbital modulation}
\IBVSabs{are found. New evidences for short lived brightenings to occur}
\IBVSabs{preferably at orbital phase ~0.35 are obtained.}

\begintext

The recurrent nova T~Coronae Borealis is classified as a symbiotic star
with one of the shortest orbital period, as well as the longest period
cataclysmic variable. H$\alpha$ observations have been obtained by
Anupama \& Prabhu (1991), Anupama (1997), and Hric et al.\ (1998).
Variability and orbital modulation of the equivalent width were
reported in these papers.

Here we present new H$\alpha$ data acquired during the last years. The
H$\alpha$ observations of T~CrB were obtained during the period
February 1993--September 2000 with the Coud\'e spectrograph of the
2.0-m RCC telescope at the Bulgarian NAO ``Rozhen'' using different CCD
detectors. The data processing has been done with the IRAF software
package. The equivalent widths (EW) of the H$\alpha$ emission lines are
measured relatively to the local continuum and are summarized in
Table~1. The typical error of the measurements is about 10 per cent.

\begin{table}
\caption{H$\alpha$ observations of T~CrB}
\begin{center}
\begin{tabular}{cr@{.}l|cr|cr|cr}
\hline
HJD & \multicolumn2{c|}{EW} & \lower6pt\hbox{HJD} & \lower6pt\hbox{EW} & \lower6pt\hbox{HJD} & \lower6pt\hbox{EW} & \lower6pt\hbox{HJD} & \lower6pt\hbox{EW} \\[-3pt]
2400000\,+ & \multicolumn2{c|}{[\AA]} & & & & & \\
\hline
49024.62 & 7&5 & 50244.29 & 16.9 & 51007.41 & 13.4 & 51247.41 & 10.3 \\
49027.62 & 6&4 & 50321.35 & 14.6 & 51028.28 & 14.9 & 51247.49 & 11.2 \\
49180.49 & 3&1 & 50564.52 & 18.4 & 51028.37 & 13.9 & 51441.30 & 5.9 \\
49225.43 & 6&8 & 50565.48 & 20.4 & 51029.28 & 14.9 & 51441.31 & 6.5 \\
49353.63 & 5&7 & 50566.36 & 21.2 & 51029.35 & 14.9 & 51632.53 & 7.3 \\
49353.65 & 5&8 & 50618.34 & 28.4 & 51030.35 & 15.1 & 51632.55 & 7.0 \\
49356.64 & 4&4 & 50705.23 & 18.2 & 51030.36 & 14.6 & 51681.31 & 4.3 \\
49376.54 & 3&7 & 50705.24 & 17.8 & 51031.26 & 14.0 & 51681.33 & 4.5 \\
49491.51 & 3&2 & 50732.19 & 30.0 & 51031.28 & 14.4 & 51717.38 & 6.5 \\
50115.55 & 5&3: & 50732.21 & 31.3 & 51091.23 & 11.5 & 51742.30 & 4.2 \\
50181.47 & 15&4 & 50881.48 & 13.8 & 51091.24 & 11.0 & 51742.32 & 4.6 \\
50182.40 & 19&6 & 50919.45 & 20.8 & 51096.25 & 9.8 & 51774.27 & 5.4 \\
50182.54 & 18&3 & 50919.50 & 21.0 & 51096.26 & 10.1 & 51774.28 & 5.5 \\
50211.39 & 12&8 & 50923.48 & 20.7 & 51184.59 & 6.4 & 51775.31 & 4.7 \\
50212.48 & 17&4 & 50969.30 & 13.4 & 51184.61 & 7.3 & 51805.27 & 2.4 \\
50242.32 & 13&3 & 51004.29 & 12.7 & 51185.63 & 8.4 & & \\
50243.31 & 14&0 & 51005.40 & 11.9 & 51185.66 & 9.5 & & \\
50244.44 & 17&4 & 51006.43 & 13.6 & 51186.58 & 7.6 & & \\
\hline
\end{tabular}
\end{center}
\end{table}

Hereafter, we will use the spectroscopic ephemeris $T_0 = {\rm
JD}~2447918.62+227\fday5687\times E$ of Fekel et al.\ (2000). The zero
epoch corresponds to a time of maximum velocity of the red giant.

The long term behaviour of the EW(H$\alpha$) is presented in Fig.~1a.
The data before JD 2448500 are from Anupama \& Prabhu (1991). The data
after it are our observations and four measurements by Miko{\l}ajewski
et al.\ (1997). A new maximum is observed sometimes between
JD~2450600--2450900 when the EW(H$\alpha$) reached values $\geq
20$~\AA. It decreased slowly to values $< 5$~\AA\ after this maximum.
This behaviour is more or less similar to variability observed around
JD~2447000. In both cases the EW reached values $\sim 30$~\AA\ and
dropped to $< 5$~\AA\ on a time scale of about 1000 days, although the
evolution in the former one seems to be slightly steeper.

The highest values in our data set are EW(H$\alpha)\geq30$~\AA. It
deserves to be noted that our highest values are at phase 0.36 of the
227\fday5687 period. Anupama \& Prabhu (1991) observed extreme values at
a close orbital phase, i.e.\ EW(H$\alpha)\geq35$~\AA\ at JD 2446860,
corresponding to orbital phase 0.35. Analyzing photographic and
visual light curves, Peel (1990) discovered a tendency for short lived
brightenings to occur at phases 0.33 and 0.20 (recalculated in terms of
the ephemeris used here). The flare like events of 1963 and 1975 (see
Palmer \& Africano 1982 and references therein) are at phases 0.55 and
0.38 respectively. All these results support the idea that short lived
eruptions occur sometimes, most probably around phase
$\sim 0.35$.

It is difficult to say what can be the reason for this short flare like
events. Most probably they are resulting from increase of mass
transfer. However, the orbit is assumed to be almost circular, with
eccentricity $e<0.02$, and perhaps this is a spurious eccentricity,
result of tidal effects (Kenyon \& Garcia 1986; Belczy{\'n}ski \&
Miko{\l}ajewska 1998). In any case it is remarkable that the majority
of the brightening events (of the EW(H$\alpha$) and the optical
magnitude) takes place around $\sim 0.35$.

\IBVSfig{18cm}{5013-f1.ps}{The EW(H$\alpha$) variability of T~CrB. {\bf
a)}~The long term behaviour. Two maxima are visible about JD 2446800
and about JD 2450700. {\bf b)}~Periodogram for the EW(H$\alpha$) using
only our data and values less than 25~\AA.  The asterisks indicate the
orbital and half orbital period. {\bf c)}~Our EW(H$\alpha$)
observations folded with the orbital period. The circles refer to ${\rm
EW}<25$~\AA, and the crosses refer to ${\rm EW}>25$~\AA}

\vskip-13pt
In the previous investigations, a modulation of the EW(H$\alpha$) with
the orbital phase has been supposed (i.e.\ Hric et al.\ 1998). We
performed periodogram analysis applying PDM (Stellingwerf 1978) and
CLEAN (Roberts et al.\ 1987) algorithms. We used the whole data set
and different subsets, with and without subtraction of fit to the data.
The fits were low order polynomial over the whole data set or over the lower
values only. Very weak traces of the orbital period are visible only when
we use our data after removing the highest values, i.e.\ using only
the points with ${\rm EW}<25$~\AA. The corresponding periodogram is
plotted in Fig.~1b.

In Fig.~1c we plotted our data folded with the orbital period of
227.5687 days. In this panel the circles represent values less than
25~\AA, and the crosses refer to ${\rm EW}>25$~\AA. As the crosses are very
different from the other measurements, we suppose that the extreme
values above 25~\AA\ are caused by short lived brightenings. If we
have in mind the circles only (i.e.\ ${\rm EW}<25$~\AA) two maxima appear to
be visible. One is at phase about 0.9--1.2 and the second at 0.6. They
are shifted relatively to the maxima detected by Hric et al.\ (1998),
even if we use the same ephemeris. It is worth noting that in Fig.~1c
the data are plotted without subtracting any fit. The fact that our
maxima are shifted relatively to those detected by Hric et al.\ (1998)
points out that the variability of the H$\alpha$ induced by the binary
rotation, is not very stable, if it exists at all. Hric et al.\ (1998)
suggested that the orbital modulation of H$\alpha$ might be a result of
the presence of two other emitting regions (which emission is
superimposed on the emission of the accretion disk) --- one identified
with the gas stream through the inner Lagrangian point ($L_1$) and the
second fed by matter rotating around the hot component but not captured
by the accretion disk. The shift of our maxima as well as the big
scatter indicate that the position of these two regions is probably
variable. This could be a result of variability of the mass transfer
rate, the angular momentum transfer rate, or/and variability of the
size of the accretion disk.

It deserves noting that the IUE observations of the integrated UV-flux
(1250--3200~\AA) during the period JD~2446000--JD~2447200 do not
exhibit considerable variations (Selvelli et al.\ 1992). In the same
time (see Fig.~1a) the EW(H$\alpha$) shows an increase from 10~\AA\ to
30~\AA\ followed by a decrease to values less than 5~\AA. Because the
integrated UV flux is a good representation of the mass accretion rate,
this points out that the variability of H$\alpha$ is probably a result
of changes in the angular momentum accretion rate and the size (and/or
the structure) of the accretion disk rather than changes in the mass
accretion rate.

High resolution observations and analysis of the H$\alpha$ emission
line profiles could throw a new light over this issue.

We acknowledge the support by the Junta de Andaluc{\'{\i}}a (Spain).

\references

Anupama, G.C., 1997, in {\it Physical Processes in Symbiotic Binaries
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Belczy{\'n}ski, K., Miko{\l}ajewska, J., 1998, {\it MNRAS}, {\bf 296}, 77

Fekel, F.C., Joyce, R.R., Hinkle, K.H., Skrutskie, M.F., 2000,
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\endreferences

\end{document}


