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% V=R transition in 66 Oph 
% IBVS version,  February 20, 1995)
\newcommand{\lam}{$\lambda$}
\newcommand{\kms}{km s$^{-1}$}
\newcommand{\beq}{\begin{equation}}
\newcommand{\eeq}{\end{equation}}
\newcommand{\Rd}{R_{\rm d}}
\newcommand{\Dv}{\Delta v}
\newcommand{\Dvd}{\Delta d_{\rm d}}
\newcommand{\D}{\displaystyle}
\newcommand{\FeII}{Fe~{\sc ii}}
\newcommand{\Ha}{H$\alpha$}
\newcommand{\Hb}{H$\beta$}
\newcommand{\VR}{$V/R$}
\input ibvs.sty


\begin{document}
\IBVShead{4174}{22 March 1995}

\centerline{\ft OBSERVATION OF A  $V=R$\/ TRANSITION IN THE Be STAR 
66 Oph\footnote{Based on observations 
obtained  at the 
European Southern Observatory, La Silla, Chile;  at the 
German-Spanish Observatory DSAZ, Calar Alto, Spain; and at the
Ond\v{r}ejov Observatory, Czech Rep.}}  
\vskip 8mm

\begintext
 
As part of our long-term high resolution spectroscopic monitoring programme of Be stars, we have repeatedly observed 
the  interesting equatorial Be star 66 Oph = HR 6712 (B2~IV-Ve, 
$v \sin i$ = 240 \kms). 
During the past two decades, its \Ha\ emission strength 
has been observed to steadily increase
from $F_\alpha$ = 3$F_{\rm c}$ in 1973 (Gray \& Marlborough 1974) to
8--10$F_{\rm c}$ in 1992--1994 (see Table 1). 
 This latter value corresponds to equivalent width
$W_\alpha \approx  50-60$\AA, 
thus presently making 66 Oph  one of the brightest Be stars at \Ha.

 
We have mainly investigated the \Ha\ and the \FeII\ $\lambda 5317$ emission lines.
Our data 
cover the epoch 1989--1994 and  have been measured  at ESO
1.4m Coud\'{e} Auxiliary Telescope (observers: Hanuschik,  Hummel),
at the 
2.2m telescope at the German-Spanish Observatory DSAZ 
on Calar Alto/Spain (observers: Hummel, Vrancken), and 
 at  
 the 2.0m Ond\v{r}ejov telescope (observer: \v{S}tefl). 
 Resolution $R = \lambda/\Delta \lambda $ has been around  
50~000, except for the Ond\v{r}ejov data (15~000).
The signal-to-noise ratio  is usually
several hundred except for some of the 
\FeII\ profiles shown here.
Profile parameters are collected in Table 1.

Our profile survey  in Fig. 1 demonstrates 
 that  in 1989--April 1993, and since June 1994 again, 66~Oph
has  shown 
asymmetric single or double-peak profiles at \Ha, and extremely asymmetric, so-called
steeple profiles in \FeII. 
Especially striking is the   inversion of asymmetry  from 
$V>R$ (in 1989) to $R>V$ (1992) back to $V>R$ (1994). 

The  steeple-type \FeII\ profile shape
and the cyclic \VR\ asymmetry are two connected phenomena, 
both being produced by a large-scale 
 density inhomogeneity.
This structure is likely to be a {\em global density
wave}, slowly precessing under the influence of the centrifugally flattened
 B star (Hanuschik et al. 1995).

In 1988,  66 Oph showed a sudden 
onset of cyclic \VR\ variability, after
at least 15 years of symmetric double-peak structure ($V \approx  R$).
Its present full \VR\  cycle time is only 5 years,
rather short if compared to other such Be stars which have typical cycle
duration of about 10 years. 
With this cycle time, the first $V=R$ transition must have occurred
in early-1991, but escaped  detection. In 1993, we have been fortunate enough to observe
 the second $V=R$\/ transition,  
both in \Ha\ and \FeII\ (see Fig. 2). 
This transition occurred in November 1993 and appears to have lasted only a
few  months, as a comparison 
of
our  data from September 1993 (\Ha: $R>V$), 
November 1993 ($R=V$), 
 and  June 1994 ($V>R$) clearly shows. 

The transition appears quite smooth, with
a gradual decrease in asymmetry in 1993 from April over September to November,
and a new increase, with opposite sign, to 1994 August.
The overall time of observed $V=R$ shape is only a few months, 
a small fraction of the full cycle time. 
Such temporal behaviour agrees well with the expectation that 
the  slope of the \VR\ evolution  is sinusoidal, with  
much longer periods of \VR\ asymmetry than with $V=R$. 
\vfill \eject

\centerline{2}
\epsfysize=15cm
\centerline{\epsffile{4174-f1.ps}}
\begin{center}
Figure 1. 
\Ha\ and \FeII\ emission line profiles in 66 Oph (left),
both on a common heliocentric velocity scale. 
The flux scale of the  \FeII\ lines is expanded by
a factor of  20.  
\end{center}

\epsfysize=6cm
\centerline{\epsffile{4174-f2.ps}}
\begin{center}
Figure 2:  Comparison of the 1993 and 1994 profiles of 66 Oph, demonstrating the $V<R
\rightarrow V>R$ transition in November 1993.
\end{center}
\vfill \eject

\centerline{3}
\vskip 5mm
\centerline{Table 1. Line parameters for \Ha\ and \FeII}
\vskip 2mm
\begin{center}
\begin{tabular}{llrlllll} \hline \noalign{\smallskip}
 Date$^a$ & \multicolumn{3}{l}{\Ha} & &\multicolumn{3}{l}{\FeII\ $\lambda$5317}\\ 
\cline{2-4}  \cline{6-8} \\ 
 & $W_\lambda$/\AA & $V/R^b$  & $F_{\rm p}/F_{\rm c}^c$ && 
  $W_\lambda$/m\AA & $V/R$  & $F_{\rm p}/F_{\rm c}$ \\ 
\noalign{\smallskip}
\hline \noalign{\smallskip}
 891001 & 45.9  & $>$ 1    & 7.78  && 680  & 3.78   & 1.189 \\
 920324 & 46.6  & 0.911    & 7.95 && 391  & 0.241  & 1.174\\
 930418 & 59.7  & 0.852    & 9.80 \\
 930909 & 59.1  & 0.944    & 9.47  && 227  & 1.59   & 1.063 \\[.5ex]
 931106 & 44.6  & 1.00     & 7.16 && 500:  & $\approx 1$  & 1.09:\\
 940624$^d$ & 50.5  & 1.10 & 7.82 \\
 940706 & 49.8  & 1.11 & 7.73\\
 940803$^e$ & 48.7 & 1.14 & 7.76\\
 940815 & 49.5  & 1.14 & 7.64\\
 940829 & 49.9  & 1.15 & 7.66 \\[1ex]
\noalign{\medskip}
\hline
\multicolumn{8}{l}{$^a$ 891001 = 1989 October 1}\\
\multicolumn{8}{l}{$^b$ $V/R = [F(V)-F_{\rm c}]/[F(R)-F_{\rm c}]$}\\
\multicolumn{8}{l}{$^c$ average intensity if two peaks exist}\\
\multicolumn{8}{l}{$^d$ averaged from 7 measurements on June 23,
 24, and 25}\\
\multicolumn{8}{l}{$^e$ averaged from 3  measurements}\\
\end{tabular}
\end{center}
\vskip 8mm


An interesting observation is that shortly before the  $V=R$\/ transition, \Ha\ and \FeII\ profiles showed slightly opposite \VR\ behaviour:
in September 1993,  $V<R\/$ in \Ha, $V > R$ in \FeII.  
This may be partly due to the fact that in 
a certain critical parameter range, 
the
superposition of
different line broadening mechanisms (causing the  winebottle-type inflections
and the profile peaks)
causes slightly inverse 
 \VR\ ratios in  \Ha\ 
and the \FeII\ line.
Alternatively, this may be indicative of a certain time lag between the 
$V=R$ transition in both lines due to very different optical depth
and therefore different contributing 
disk regions.
Thus these measurements  contain valuable information for the 
modelling of the density wave.

 
A very pronounced decrease in equivalent width occurred in November
1993 (by 25 \%),
after the star had shown a stable value of 59 \AA\ for almost half a year
before.
Half a year later,  $W_\alpha$  was again observed at  higher values,
remaining constant thereafter for at least two months. 
Such rather strong variability seems to be  uncommon in this star. 
However, we are not fully sure that the sudden decrease is not 
a mere chance coincidence.
If it is physically related to the $V=R$ transition,
then this observation may be interpreted
as result of the relatively small velocity gradient (averaged across the
whole emitting disk) at the moment of symmetry ($V=R$), as compared
to the situation shortly before and after the transition when the
velocity gradient becomes larger again.  
\vfill \eject

\centerline{4}
\vskip 1mm

We strongly encourage other observers to continue to monitor this interesting
Be star at high spectral resolution, in order to follow up its \VR\ behaviour
and to furthermore document its emission strength variability pattern.  

\vskip 1cm

\beginauthors
\authorsrightcol
R.W. HANUSCHIK$^1$, 
S. \v{S}TEFL$^2$, 
W. HUMMEL$^3$, 
M. VRANCKEN$^3$
\endauthors
\vskip2mm

$^1$Astronomisches Institut, Ruhr-Universit\"at, 
Postfach 10~21~48, D-44780 Bochum, 

\hspace*{2mm}Germany

$^2$Astronomical Institute, Czech Acad. of Sciences, CZ-251 65 Ond\v{r}ejov, 
Czech Republic

$^3$Astrofysisch Instituut, Vrije Universiteit Brussel, Pleinlaan 2, B-1050 Brussel,

\hspace*{2mm}Belgium


\references

Gray, D.F., Marlborough, J.M., 1974, {\it ApJS}, {\bf 27}, 121

Hanuschik, R.W., Hummel, W., Dietle, O., Sutorius, E., 1995, {\it A\&A} (in press)

\end{document}

 
