\documentclass[journal]{rmaa}


%\bibpunct{(}{)}{;}{a}{}{,}
\renewcommand{\mag}[1]{^{\rm m}\!\!\!#1\,}


\title{Spectral Variability of Romano's Star} 

\author{O. Maryeva,\altaffilmark{1, 2} and P. Abolmasov\altaffilmark{3}}

\altaffiltext{1}{Stavropol State University, Russia.}
\altaffiltext{2}{Special Astrophysical Observatory, Russia.}
\altaffiltext{3}{Sternberg Astronomical Institute, Russia.}              


\shortauthor{Maryeva \& Abolmasov}

\shorttitle{Spectral Variability of V532}


\listofauthors{O.~Maryeva \& P.~Abolmasov}

\indexauthor{Maryeva, O.}
\indexauthor{Abolmasov, P.}


\fulladdresses{
\item O. Maryeva: Stavropol State University, Pushkina Str., 1, Stavropol, Russia 355009, Russia and Special Astrophysical Observatory, Nizhnij 
Arkhyz, Zelenchukskij, Karachai-Cirkassian Republic, Russia 369167 (olga.maryeva@gmail.com).
\item P. Abolmasov: Sternberg Astronomical Institute, Moscow State University, Universitetsky pr., 13, Moscow 119992, Russia (pavel.abolmasov@gmail.com).}

\SetYear{2010}
\SetVolume{46}
\ReceivedDate{2010 February 3} 
\AcceptedDate{2010 June 10} 


\resumen{Combinamos observaciones espectrosc\'opicas de archivo de la estrella LBV  V532 (conocida como Estrella de Romano) con fotometr\'ia 
en la banda $B$.  Los datos espectrosc\'opicos cubren un intervalo de 15 a\~nos, desde 1992 hasta 2007. Mostramos que en m\'aximo de brillo 
el objeto se comporta como una supergigante con l\'ineas de emisi\'on, mientras que en m\'inimo V532 se mueve a lo largo de la secuencia de 
las estrellas WN tard\'ias. En este sentido, la estrella tiene un comportamiento similar al de las bien conocidas variables luminosas azules 
(LBV:  luminous blue variables) AG Car y R 127, pero es un poco m\'as caliente en m\'inimo. Identificamos unas 100 l\'ineas espectrales en 
el intervalo 3700--7300~\AA. Nuestra espectroscop\'ia para este objeto es la m\'as completa en la actualidad. Obtenemos la velocidad del viento 
usando las l\'ineas del triplete de He~I ($300 \pm 30$~km~s$^{-1}$). Tambi\'en estimamos los par\'ametros f\'isicos de la nebulosidad que rodea 
a V532.\smallskip}

\abstract{We combine archival spectral observations of the LBV star V532 (Romano's star) together with existing photometric data in the $B$ band. 
Spectroscopic data cover 15 years of observations (from 1992 to 2007). We show that the object in maximum of brightness behaves as an emission 
line supergiant while in minimum V532 moves along the sequence of late WN stars. In this sense, the object behaves similarly to the well-known 
Luminous Blue Variable (LBV) stars AG Car and R127, but is somewhat hotter at the minima. We identify about 100 spectral lines in the $3700-7300$~\AA\ 
wavelength range. As of today, our spectroscopy is the most comprehensive for this object. The velocity of the wind is derived using the He~I 
triplet lines ($360\pm30$~km~s$^{-1}$). Physical parameters of the nebula around V532 are estimated.\medskip}


\addkeyword{galaxies: individual (M33)}
\addkeyword{stars: individual (Romano's star)} 
\addkeyword{stars: Wolf-Rayet}
\addkeyword{supergiants}



\begin{document}

\maketitle


\section{Introduction} 

Luminous Blue Variables (LBVs) are a class of rare astrophysical objects introduced by \citet{conti84} and remain a subject of intense interest. 
LBVs are broadly accepted to be very massive and energetic stars emitting close to the Eddington limit, evolving from Of toward Wolf-Rayet stars. 
However, contemporary evolutionary models do not answer the question about the exact relations between LBV, nitrogen-rich Wolf-Rayet (WN) and 
hydrogen-rich WN (WNH) stars. Numerical simulations show that the objects may pass the LBV evolutionary stage either before or after the WNH stage 
\citep{SmithConti}. Currently, only 35 LBV and LBV candidates are known in our Galaxy (Clark, Larionov, \& Arkharov 2005). Studying LBVs in nearby 
galaxies is very important for understanding stellar evolution and the evolution of the interstellar medium perturbed and contaminated by massive 
stars at various evolutionary stages.


%%%%% TABLA 1 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\begin{table*}[!t]
\vspace{0.1cm}
\centering
\caption{Observational log for MPFS data. S/N is signal-to-noise ratio\\ per resolution element}
\label{tab:obstabmpfs}
\small
\begin{tabular}{cccccc}
\toprule
Date       & Exposure time (s) & Seeing ($''$) & S/N & Spectral standard star & Spectral range (\AA)  \\
\midrule
05 10 2002 &  ~900             & 3.8           & 16  & BD25d4655             & 4250--6700   \\
13 11 2004 &  4200             & 1.5           & 20  & HZ44                  & 4000--7000   \\
17 01 2005 &  3600             & 1.5           & 18  & G248                  & 4000--7000   \\
\bottomrule         
\end{tabular}
\vspace{0.4cm}
\end{table*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


The object V532 ($\alpha=01^{h} 35^{m} 09.^{s}71$, $\delta=+30^{\circ}41'57\rlap.{''}1$) is located in the outer spiral arm of the M33 galaxy. Giuliano 
Romano was the first to recover its light curve  \citep{romano} and to find irregular magnitude variations between 16$\mag{.}7$ and $18\mag{.}1$. 
Romano classified V532 as a variable of the Hubble-Sandage type by the shape of the light curve and its colour index. \citet{HumphreysDavidson} 
on the basis of its light variations classified the star as an LBV candidate. Two maxima of brightness were detected during the last half of the 
century \citep{kurtev}. The first was observed around 1970 and the second in the early 1990s \citep{kurtev}. Long-term variability has an amplitude 
of about 1$\mag$\,\, and seems to be superimposed on an even stronger downward trend. Short-timescale variability with amplitude $\sim 0\mag{.}5$ 
was discovered in addition to the longer-timescale variability \citep{kurtev,olgaalla}. Such photometric behavior is typical for an LBV star.

The first optical spectrum was obtained by T. Szeifert (1996) at the 3.5~m Calar Alto telescope in 1992. In this spectrum, ``few metal lines are visible, 
although a late B spectral type is most likely (faint HeI)''. This spectrum, obtained with the TWIN spectrograph of the Calar Alto telescope, is 
unique in being obtained in a profound flare state. It is drastically different from the hot spectra observed at the minima of brightness (see below). 

Another spectrum was obtained by O. Sholukhova at the 6 m telescope of the Special Astrophysical Observatory (SAO) of the Russian Academy of Sciences 
(RAS) in 1994 \citep{olga}. Another spectrum was obtained at SAO with the Multi-Pupil Fiber Spectrograph (MPFS) in September 1998. This spectrum 
was classified as WN10--WN11 \citet{fabrika}. \citet{polcaro} estimated the bolometric absolute magnitude of the object as $M_\mathrm{bol} \simeq -10\mag{.}4$, 
using a bolometric correction ``of at least --3~mag'' and the distance modulus $m-M=24\mag{.}8$. They classify V532 as an LBV because the object fulfills 
all the criteria of \citet{HumphreysDavidson}. Using five spectra obtained in 2003--2006, \citet{viotti2006,viotti2007} find an anti-correlation 
between equivalent widths of the Wolf-Rayet blue bump at 4630--4686~\AA\ and visual luminosity.

Comparing  the spectra  published by \citet{szeifert}, \citet{fabrika} and \citet{viotti2007}, we find that the object changes its spectral properties 
significantly. 

Here we combine the spectral observations (both new and already published) with the light curve of the object to trace the spectral variability of the 
star. We describe the data and data reduction process in the next section. Results are presented in \S~\ref{sec:res} and discussed in \S~\ref{sec:disc}. 
In \S~\ref{sec:con}, we present the conclusions.


%%%%% TABLA 2  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\begin{table*}[!t]
\centering
\caption{Observational log for the SCORPIO data. $\alpha$ is seeing, PA is position angle,\\ S/N is signal-to-noise ratio}
\label{tab:obstabscor}
\setlength{\tabcolsep}{1.2\tabcolsep}
\tablecols{8}
\small
\begin{tabular}{ccccccccc}
\toprule
Date & Exposure  & Grism & Spectral    & $\delta \lambda$ (\AA) & S/N  & $\alpha ('')$ & Spectral      & PA ($^{\circ}$) \\
     & time (s)  &       & range (\AA) &                        &      &               & standard star &  \\
\midrule
06 02 2005 & ~600  &           &             &     & ~8  & 1.7  & G248~~~~  & $-$136 \\
30 08 2005 & 1200  & VPHG550G~ & 3500--7200  & 10  & 30  & 1.9  & G191-B2B~ & ~~210  \\
08 11 2005 & 3300  &           &             &     & 45  & 1.9  & BD25d4655 & ~~145  \\
\noalign{\smallskip}
03 08 2006 & 1500  &           &             &     & 20  & 2~~  & GD248~~~~ & ~~200  \\
10 08 2007 & 1800  & VPHG1200G & 4000--5700  & ~5  & 24  & 2~~  & BD33d2642 & ~~252  \\
05 10 2007 & 2700  &           &             &     & 40  & 1.1  & BD25d4655 & $-$141 \\
\noalign{\smallskip}
08 01 2008 & 1800  & VPHG1200R & 5700--7500  & ~5  & 20  & 2.1  & BD25d4655 & ~~~48  \\
\noalign{\smallskip}
10 01 2008 & 1800  & VPHG1200G & 4000--5700  & ~5  & 20  & 1.4  & BD25d4655 & ~~~18  \\
\bottomrule
\end{tabular}
\vspace{0.45cm}
\end{table*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


\section{Observations and Data Reduction}
\label{sec:obs}

In this work, we use archival data from the 6m SAO telescope\footnote{Available via ASPID database \url{http://alcor.sao.ru/db/aspid/}.} and 
the SUBARU telescope, which is operated by the National Astronomical Observatory of Japan. The 6 m telescope data were obtained with the Multi 
Pupil Fiber Spectrograph (MPFS) (Afanasiev, Dodonov, \& Moiseev 2001) and with the SCORPIO multi-mode focal reducer in the long-slit mode 
\citep{scorpio}. The data from SUBARU were obtained with the Faint Object Camera (FOCAS) \citep{focas} in the Cassegrain focus. Uncertainties 
are everywhere dominated by statistical Poissonian noise (readout noise and round-off errors are significantly smaller).

\subsection{MPFS Data}

MPFS obtains simultaneously the spectra from 240 ($16\times 15$) or 256 ($16\times 16$) spatial sampling elements arranged in the form of 
a rectangular array of square lenses. In 2002 and 2004--2005, the data were obtained in the 240-element and 256-element configurations, 
respectively. Spatial sampling of $1'' \times 1''$ was used. Light from individual sampling elements is collected by microlenses and 
transmitted by means of optical fibers reorganized in the form of a pseudo-slit toward the spectral camera. Grating \#~4 (600~\AA~pix$^{-1}$) 
providing a spectral resolution of about 6~\AA\ was used for all the observations. Detectors CCD TK 1024 ($1024\times 1024$ pixels) and 
EEV42-40 ($2048\times 2048$ pixels) were used in 2002 and 2004--2005, respectively. A sky background spectrum at a distance of $4''$ away 
from the object is taken simultaneously with the object by 17 (for $16 \times 16$ field size) or 16 (for $16 \times 15$ configuration) 
additional fibres. We summarize the relevant information on the MPFS data in Table~\ref{tab:obstabmpfs}. 

The data reduction system was written in the IDL environment and makes use of procedures written by V. Afanasiev, A. Moiseev and P. Abolmasov. 
The reduction process consists of the standard steps for panoramic data reduction (see for example S\'anchez 2006): bias subtraction, flat-fielding, 
removal of cosmic-ray hits, extraction of the individual spectra from the frames and their wavelength calibration using the spectrum of a He-Ne-Ar 
lamp. At every wavelength, we calculate the median sky level using the offset fibres and subtract from the spectra of the field. Spectra 
of spectrophotometric standard stars were used for absolute flux calibration. 

Three emission-line spectra obtained with MPFS in 2002--2005 were used. We estimate the signal-to-noise ratio in continuum as 10--30 per resolution 
element (see Table~\ref{tab:obstabmpfs} for details). Integral spectra were extracted in an annular aperture $2''$ in radius. Up to the instrumental 
resolution, the object is point-like. Note that, unlike long-slit spectrographs, MPFS being free from slit losses and absolute flux calibration, is much 
more reliable. 


%%%%% FIGURA 1

\begin{figure}[!t]
\vspace{0.08cm}
\centering
\includegraphics[width=0.99\columnwidth]{omaryevafig1.eps}
\caption{$B$-band light curve of V532. Arrows indicate when spectroscopic observations were made.  
Data from \citep{viotti2007} are shown by diamonds.}
\label{fig:lightcurve}
\vspace{0.3cm}
\end{figure}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


\vspace{0.2cm}

\subsection{SCORPIO Data}
\label{sec:scorpio}

Observational data from SCORPIO are summarized in Table~\ref{tab:obstabscor}. All the spectra were reduced using the {\tt ScoRe} package for 
long-slit data reduction, written in IDL\footnote{Available at \url{http://narod.ru/disk/24196358000/Score_v1.2.tar.html}.}. CCD frames were 
de-biased, cosmic particle hits were removed from all types of exposures except bias. Save for the single 10 minute long exposure obtained 
in February 2005, several object exposures are present for every day, which simplifies the removal of cosmic hits. Then we divide object and 
spectral standard exposures over the normalized mean flat-field frame. After wavelength calibration using He-Ar-Ne lamps and night sky OI 
($\lambda~5577$ or $\lambda~6300$ depending on grism) emission lines, the CCD data were flux-calibrated using standard stars from Oke's (1990) 
spectrophotometric standard list (see Table~\ref{tab:obstabscor}). Spectra were extracted by fitting the profiles of slices across dispersion 
with a Gaussian function.  

In total, eight spectra were obtained with SCORPIO in 2005--2008 (dates of observations are shown by upward arrows in Figure~\ref{fig:lightcurve}). 
The spectral resolution is 10~\AA\ (for VPHG550G grism) and 5~\AA\ (for VPHG1200G and VPHG1200R grisms), the signal-to-noise ratios of the continuum 
vary from about 10 to about 50 for larger total exposures.  


%%%%% FIGURA 2 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\begin{figure}[!t]
\centering
\includegraphics[width=0.99\columnwidth]{omaryevafig2.eps}
\caption{H$\alpha$ line profile in the Calar Alto spectrum obtained in 1992. The spectrum 
is normalized by the local continuum level.}
\label{fig:twinhalpha}
\vspace{0.2cm}
\end{figure}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


\subsection{Data from SUBARU}

One exposure 1200s in length was obtained with the SUBARU telescope on October 2007. The VPHG450 grism was used providing the spectral range of 
3750--5250~\AA. A slit width of 0$\rlap.{''}5$ implies a spectral resolution of about 1.7~\AA\ (currently, it is the best resolution achieved for 
this object) in the extracted spectrum. Data were reduced using IDL-based software. The CCD frames were bias-subtracted and  flat-fielded. 
We used the Th-Ar arc spectrum for wavelength calibration. Star BD40d4032 from the SUBARU spectrophotometric standard list  was used to calibrate 
the stellar spectra. We extracted the spectrum in the same way we did for SCORPIO (see above \S~\ref{sec:scorpio}).

If we degrade the spectral resolution to 5~\AA, the spectrum becomes practically identical to the spectrum obtained with SCORPIO on the same 
date. Though the spectral shapes are similar, the spectra differ in normalization (by about a factor of three) connected to slit losses.


%%%%% FIGURA 3

\begin{figure*}[!t]
\centering
\includegraphics[width=0.8\textwidth]{omaryevafig3.eps}
\caption{Optical spectrum evolution in the blue spectral range (4000--5200~\AA). Spectra 
are normalized by the local continuum level and vertically shifted for clarity.}
\label{fig:spevolution}
\vspace{0.4cm}
\end{figure*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


\subsection{Photometric Data}

The light curve (see Figure~\ref{fig:lightcurve}) was provided by Vitalij Goranskij and consists of the CCD data obtained by Goranskij and 
Zharova with the 1 m SAO telescope and two instruments of the Crimean laboratory of the Sternberg Astronomical Institute (SAI), and photographic 
plates from the SAI collection. CCD data were reduced with MaxIm DL software\footnote{\url{http://www.cyanogen.com/maxim_main.php}.}. The joint 
light curve will be published by A.~V. Zharova et al. (in preparation) in a separate paper, all the details of the data reduction process will 
be given there. Photographic B-band magnitudes were identified with the B-band magnitudes obtained by CCD observations. CCD data uncertainties 
are of the order of 0.05~mag, plate data have larger errors depending on the source brightness, usually of the order 0.1--0.2~mag.

The joint curve contains the data of \citet{viotti2007} but complements them with much larger photometric material and primarily allows to trace 
the recent evolution of the object.


\section{Results}
\label{sec:res}

\subsection{Spectral Evolution} 

The first optical spectrum of V532  was obtained during the rise of brightness, when the object was 16$\mag{.}$4 in the $V$ band. This 
spectrum covers two spectral ranges $\lambda\lambda~4400-5000$ and $\lambda\lambda~5800-6750$, as it was obtained with the two-channel 
TWIN spectrograph\footnote{A comprehensive description is present at \url{http://www.caha.es/CAHA/Instruments/TWIN/HTML/twin.html}.}. 
In this spectrum, two strong lines are present, H$\alpha$ (shown in Figure~\ref{fig:twinhalpha}) and H$\beta$, having complex profiles 
consisting of a narrow line with P~Cyg profile and broad wings. FWHM (Full Width at Half Maximum intensity) of the wings is $\sim 15$~\AA\ 
for H$\beta$ and $\sim 20$~\AA\ for H$\alpha$. Such wings were first found for LBVs in the spectrum of P Cygni by \citet{bernat}. These 
wings are explained by scattering of line photons by free electrons in the stellar wind. Similar line profiles were observed for the LBV 
stars R127 and AG~Car during the initial rise to maximum \citep{stahl1983,stahl2001}. 


%%%%%% TABLA 3 (ERA LA 4) %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\begin{table}[!t]
\vspace{0.1cm}
\centering
\tablecols{3}
%\setlength{\tabnotewidth}{1\textwidth}
\setlength{\tabcolsep}{1.5\tabcolsep}
\caption{Spectral classes identified\\ via the scheme of Smith,\\ Crowther, \& Prinja (1994)}
\label{tab:spsmith}
\small
\begin{tabular}{llc} 
\toprule
\multicolumn{1}{c}{B, mag} & \multicolumn{1}{c}{Spectral} & Date \\
                           & \multicolumn{1}{c}{subtype}  &      \\
\midrule 
17.5    & WN10.5   & 2002/10/05  \\
16.9    & WN11     & 2004/11/13  \\  
17.1    & WN11     & 2005/01/17  \\ 
17.15   & WN11     & 2005/02/06  \\ 
17.3    & WN10     & 2005/08/30  \\    
17.6    & WN9      & 2005/11/08  \\
18.3    & WN8      & 2006/08/03  \\
18.4    & WN8      & 2007/08/10  \\ 
18.5    & WN8      & 2007/10/05--08 \\  
        & WN8      & 2008/01/08--10 \\
\bottomrule
\end{tabular}
\vspace{0.2cm}
\end{table}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


Emission lines of He~I~$\lambda~5876$ and Si~II~$\lambda~5957.612$, 5978.970, 6347.091, 6371.359 are also present in the red spectrum. 
\citet{szeifert} mentions these SiII lines as ``weak metal emissions''.

Figure~\ref{fig:spevolution} shows all the spectra of V532 in the blue range (4000--5500~\AA) analysed in this article, obtained at different 
spectral resolutions. All the spectra are normalised to continuum level in a uniform way. In order to get a reasonable normalisation for our 
spectra, we chose several wavelength intervals practically free from Wolf-Rayet emissions (4250--4270, 5100--5200, 5520--5620, 5750--5800 and 
6950--6970~\AA) and reconstructed the continuum using a second-order polynomial. Below, for spectral classification we use characteristic 
equivalent width ratios, which makes our results practically independent of spectral resolution. 

The spectral appearance of the spectra of V532 obtained in 2002--2008 resembles that of late WN stars. It allows us to apply a classification scheme 
used for WN stars, bearing in mind that abundances of individual elements may differ, but physical conditions are similar. All the spectra obtained 
between 2002 and 2008 were classified using the classification of Smith, Crowther, \& Prinja (1994) for WN6-11 stars based primarily on relative 
strengths of N~V~$\lambda\lambda~4604-20$, N~IV~$\lambda~4058$, N~III~$\lambda\lambda~4634-41$ and N~II~$\lambda~3995$ emission lines. The method has 
low dependence on elemental abundances, though only helium and nitrogen lines (preferably, ratios of the lines of one element) are used. The results 
of the spectral classification are given in Table~\ref{tab:spsmith}.


%%%%% FIGURA 4 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\begin{figure*}[!t]
\centering
\includegraphics[clip,width=0.99\textwidth]{omaryevafig4.eps}
\caption{Left panel: V532 location on the equivalent width diagram of He~I~$\lambda~5876$ versus He~II~$\lambda~4686$ 
for different dates of observation. The points are connected with dotted lines in chronological order. Right panel: 
V532 location in October 2007 (FOCAS data) on the equivalent width versus FWHM diagram for the He~II~$\lambda~4686$ 
line. In both graphs, known Galactic (open symbols) and LMC (filled) WN stars are shown for comparison: WN8 by 
circles, WN9 by squares, and WN10 by triangles. Data for these objects were taken from \citet{smith}.}
\label{fig:spclass}
\vspace{0.4cm}
\end{figure*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


%%%%% TABLA 4 (ERA LA 3) %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%5

\begin{table*}[!t]
\vspace{0.1cm}
\centering
\caption{List of emission lines detected in the FOCAS and SCORPIO\\ (January 2008) spectra of V532\tabnotemark{a}} 
\label{tab:linelist}
\tablecols{9}
\setlength{\tabnotewidth}{0.99\textwidth}
\small
\begin{tabular}{clcccclcc}
\toprule
$\lambda$, \AA\ & Ion & EW emission & EW absorption  &  & $\lambda$, \AA\ & Ion & EW emission & EW absorption \\
                &     & \AA\        & \AA\           &  &                 &     & \AA\        & \AA\ \\
\cmidrule{1-4}
\cmidrule{6-9}
3770.60  &  H11+HeII          &              &              &   &   4613.90  &  NII             &               &  \\
3797.90  &  H10+HeII          &              &              &   &   4621.40  &  NII             &               &  \\ 
3819.76  &  HeI               &              &              &   &   4630.54  &  NII             &               &  \\
3835.39  &  H9+HeII           &              &              &   &   4634.00  &  NIII            & $6.2~\pm1.0$  &  \\
3871.82  &  HeI               &              &              &   &   4640.64  &  NIII            & $7.4~\pm0.5$  &  \\
3889.05  &  H8+HeI            & $11~\pm1~~$  & $3.7\pm1.3$  &   &   4643.09  &  NII             &               &  \\
3964.73  &  HeI               &              &              &   &   4650.16  &  C III           & $5.0~\pm0.5$  &  \\
3970.08  &  H$\epsilon$+HeII  &              &              &   &   4658.10  &  [FeIII]         & $2.5~\pm1.0$  &  \\
3994.99  &  NII               &              &              &   &   4685.81  &  HeII            & $15.2\pm0.2$  &  \\
4009.00  &  HeI               &              &              &   &   4701.50  &  [FeIII]         & $2.9~\pm1.0$  &  \\
4025.60  &  HeI+HeII          & $2.5\pm0.5$  & $1.2\pm0.2$  &   &   4713.26  &  HeI             & $1.6~\pm0.2$  &  \\
4088.90  &  SiIV              &              &              &   &   4861.33  &  H$\beta$        & $22.5\pm1.5$  &  \\
4097.31  &  N III             &              &              &   &   4921.93  &  HeI             & $4.3~\pm0.2$  &  \\
4101.74  &  H$\delta$+HeII    & $7.6\pm1.0$  &              &   &   4958.91  &  [O III]         & $1.3~\pm0.2$  &  \\
4103.40  &  NIII              &              &              &   &   5006.84  &  [O III]         & $3.9~\pm0.6$  &  \\
4116.10  &  SiIV              &              &              &   &   5015.67  &  HeI             & $2.6~\pm0.3$  &  \\
4120.99  &  HeI               &              &              &   &   5411.50  &  HeII            & $6.2~\pm0.2$  &  \\
4143.76  &  HeI               &              &              &   &   5666.60  &  NII             &               &  \\
4199.80  &  HeII              &              &              &   &   5676.02  &  NII             &               &  \\
4236.93  &  NII               &              &              &   &   5679.56  &  NII             &               &  \\
4241.79  &  NII               &              &              &   &   5686.21  &  NII             &               &  \\
4241.79  &  NII               &              &              &   &   5710.76  &  NII             &               &  \\
4340.47  &  H$\gamma$+HeII    & $8.0\pm0.3$  &              &   &   5875.79  &  HeI             & $25.5\pm0.7$  &  \\
4387.93  &  HeI               & $1.4\pm1.0$  &              &   &   6548.00  &  [N II]          &               &  \\
4471.69  &  HeI               & $5.0\pm0.3$  & $1.7\pm0.2$  &   &   6562.82  &  H$\alpha$+HeII  & $107~\pm3~~$  &  \\
4481.13  &  MgII              &              &              &   &   6583.00  &  [N II]          &               &  \\
4510.90  &  NIII              &              &              &   &   6678.15  &  HeI             & $18~~\pm3~~$  &  \\
4514.90  &  NIII              &              &              &   &   6683.20  &  HeII            &               &  \\
4518.20  &  NIII              &              &              &   &   7065.44  &  HeI             & $18~~\pm3~~$  &  \\
4523.60  &  NIII              &              &              &   &   7135.73  &  ArIII           & $3.2~\pm0.3$  &  \\
4530.80  &  NIII              &              &              &   &   7281.35  &  HeI             &               &  \\
4534.60  &  NIII              &              &              &   &            &                  &               &  \\
4541.60  &  HeII              &              &              &   &            &                  &               &  \\
4547.30  &  NIII              &              &              &   &            &                  &               &  \\
4601.50  &  NII               &              &              &   &            &                  &               &  \\
4607.20  &  NII               &              &              &   &            &                  &               &  \\
\bottomrule
\tabnotetext{a}{For higher signal-to-noise ratios, equivalent widths (EW) are given. For lines with P~Cyg profiles, 
we give both emission and absorption component EWs.}
\end{tabular}
\vspace{0.25cm}
\end{table*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


The light curve exhibits a local maximum in 2004 and early 2005. In the spectra obtained in this period, one may observe strong emission lines 
of hydrogen and neutral helium. The spectral appearance of V532 shows strong similarities with a spectrum of a WN11 star. He~I~$\lambda~4713$ 
line is stronger than the N~II blend, and He~II~$\lambda~4686$ is absent. In the spectrum obtained in 2002, the N~II~$\lambda~3995$ line is outside 
the spectral range; therefore we carry out the classification comparing N~III~$\lambda\lambda~4634-41$ and N~II~$\lambda\lambda~4601-43$ blends having 
similar intensity to the N~II~$\lambda~3995$ emission. N~III~$\lambda\lambda~4634-41$ lines appear, but are weaker than N~II~$\lambda\lambda~4601-43$. 
He~II~$\lambda~4686$ line is approximately as bright as the N~II~$\lambda\lambda~4634-41$ emission; hence we classify the object as WN10.5. 

Starting from the middle of 2005, Romano's star weakens in the optical range; its visible magnitude falls from 17 to 18$\mag{.}$8. During this 
period its spectrum evolves from WN10 (August 2005) through WN9 (November 2005) to WN8 (August 2007) by three spectral sub-classes. In October 
2007, He~II~$\lambda~4686$ line is stronger than N~III~$\lambda\lambda~4634-41$. However, the N~IV lines do not appear in the spectrum. We also classify 
the spectrum as WN8.

The evolution of V532 on the equivalent width diagram of He~I~$\lambda~5876$ versus He~II~$\lambda~4686$ is shown in the left panel of Figure~\ref{fig:spclass}. 
Locations of well-proven Galactic and Large Magellanic Cloud (LMC) Wolf-Rayet stars R84 (WN9), Sk-60~40 (WN10), LBV HDE~269582 (WN10) and 
some other objects are shown for comparison with our object. The figure shows that in the period between 2002 and 2005, when the star 
brightened by about one magnitude, its spectrum changed from WN9.5 to WN11. Since 2005, the spectrum has changed smoothly according to the 
sequence established by Smith, Shara, \& Moffat (1996).


%%%%% FIGURA 5 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\begin{figure*}[!t]
\centering
\includegraphics[width=0.82\textwidth]{omaryevafig5.eps}
\caption{Optical spectrum of V532 obtained with FOCAS in the range 3800--5100~\AA\AA.}
\label{fig:subaruspect}
\vspace{0.3cm}
\end{figure*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


We used a quantitative chemistry-independent criterion based on the FWHM of the He~II~$\lambda~4686$ line for an alternative spectral classification. 
In Figure~\ref{fig:spclass}, we show the location of V532 on the diagram of the equivalent width of He~II~$\lambda~4686$ versus the FWHM of this 
line. We have measured equivalent width and FWHM of the He~II line in the FOCAS spectrum. For this, we approximate the Wolf-Rayet blue bump by 7 
Gaussians. The position of V532 is fully consistent with its Galactic and LMC WN9 analogues. The spectral class defined from the diagram is 
consistent with that determined from the relative strengths of N~II, N~III and He~I with an accuracy of about one subclass. 

Table~\ref{tab:linelist} shows the lines detected in the spectra obtained with FOCAS (October 2007) and with SCORPIO (January 2008, with 
grisms 1200G and 1200R). Three spectra are used in order to cover the maximal wavelength range at maximal possible spectral resolution. 
The spectrum of V532 does not change noticeably from October 2007 to January 2008. Equivalent widths of the principal lines (absorptions, 
emissions and P~Cyg) are given with errors of approximation. Uncertainties due to the choice of the continuum level are smaller than the 
errors of approximation. Lines with P Cygni profiles were resolved by fitting with the models described in the next subsection. 


\subsection{Line Profiles and Terminal Wind Velocity} 

In Figure~\ref{fig:subaruspect} we show the FOCAS spectrum in the optical blue range $3800-5100$~\AA. The Wolf-Rayet blue bump (consisting primarily 
of N~III~$\lambda\lambda~4634, 4640$, C~III~$\lambda~4650$, He~II~$\lambda~4686$ and He~I~$\lambda~4713$) is clearly seen in this spectrum. We suppose 
that the lines near 4658 and 4701~\AA\ are nebular lines of [Fe~III]. They cannot belong to C~IV since C~IV~$\lambda\lambda~5812, 5801$ lines are not 
present in the SCORPIO spectrum obtained simultaneously in October 2007.


%%%%% FIGURA 6 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\begin{figure*}[!t]
\centering
\includegraphics[width=0.995\textwidth]{omaryevafig6.eps}
\caption{Top panel: profiles of triplet He~I~$\lambda~3889,~4025,~4471$ lines (from left to right). 
Two-component model fits are shown (see text for details). Bottom panel: profiles of singlet 
He~I~$\lambda~3965,~4922,~5016$ lines. FOCAS data.}
\label{fig:singlet}
\vspace{0.4cm}
\end{figure*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


Analyzing the FOCAS spectrum of V532, we found that triplet and singlet lines of HeI have different profiles (see Figure~\ref{fig:singlet}). 
Triplet lines of He~I~($\lambda~3889,~4025,~4471$) show strong P~Cyg profiles, while singlet lines ($\lambda~3965,~4922,~5016$) have flat-topped 
profiles. The widths of these lines correspond to a velocity span of about 100~km~s$^{-1}$. 

We used P Cyg profiles of triplet HeI lines to estimate the terminal wind velocity $v_{\infty}$. Line profiles were fitted with a sum of two 
Gaussians (one representing the absorption, the other the emission component). We suppose that the widths of the emission and absorption components 
are equal to the instrumental profile width of 1.8~\AA\ which is almost independent of wavelength. The instrumental profile width is determined 
using wavelength calibration spectrum lines of similar wavelengths. The terminal wind velocity is estimated by the velocity shift between the 
Gaussian centers. Even for the emission component of He~I lines, we do not detect any Doppler broadening. This may mean that the profiles are 
not true P~Cyg but wind blueshifted absorptions plus nebular emissions. This possibility does not however alter the wind velocity estimates. 

Wind velocities for all the He~I lines with P~Cyg profiles are equal within the statistical errors. The mean wind velocity for the three triplet 
lines is $360\pm 30$~km~s$^{-1}$. This value is consistent with the terminal wind velocities for late WN  (Crowther, Hillier, \& Smith 1995; 
Smith, Crowther, \& Willis 1995) which we give in Table~\ref{tab:windvelocity} for comparison. 

It is interesting to compare these profiles with the He~II~$\lambda~5412$ line in the SCORPIO spectrum obtained in October 2007. The spectrum has 
very high signal-to-noise ratio which allows reasonable terminal wind velocity estimates in spite of the significantly worse spectral resolution. 
The line has a classical low-optical-depth P~Cyg profile: broad emission and blueshifted narrow absorption. We approximate this line by a classical 
P~Cyg model profile (flat-topped emission with a narrow blueshifted absorption component) convolved with an instrumental point spread function that 
we consider Gaussian (the instrumental profile width is $\sim 5.5$~\AA). Similar line profiles are produced by optically thin envelopes expanding 
at a constant velocity. Its fitting yields the wind speed of $200\pm 15$~km~s$^{-1}$. The expansion velocity indicates that the line is formed in the 
hotter inner parts of the wind that move with smaller outward velocities. Among Pickering emissions, we also detect $\lambda~4541.6$ and $\lambda~4199.8$ 
having similar profiles in the FOCAS spectrum, but we lack signal-to-noise to obtain reliable expansion velocity estimates for these lines. 


%%%%% TABLA 5 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\begin{table*}[!t]
\vspace{0.1cm}
\centering
\caption{Wind velocities estimated through optical\\ HeI lines and ambient oxygen abundances}
\label{tab:windvelocity}
\setlength{\tabnotewidth}{0.75\textwidth}
\tablecols{6}
\small
\begin{tabular}{lccccc} 
\toprule
Star  & Galaxy  & WN       & ${\rm v_{\infty}}$  & $12+\log{\rm O/H}$  & HII \\
      &         & subtype  & ${\rm km~s^{-1}}$  & abundance           & region \\
\midrule
V532 (during                              & M33  & ~9  & 360  & $8.26\pm0.08$    &       \\
minimum epoch 2007)                       &      &     &      &                  &       \\
MCA1-B\tabnotemark{a}                     & M33  & ~9  & 420  & $8.315\pm0.061$  & NGC588\tabnotemark{e} \\
B517\tabnotemark{b}                       & M33  & 11  & 275  & $8.334\pm0.083$  & MA2\tabnotemark{e}  \\
Sk-66$^{\circ} 40$\tabnotemark{a}          & LMC  & 10  & 300  &                  &      \\
R84\tabnotemark{a}                        & LMC  & ~9  & 400  &                  &       \\
BE381\tabnotemark{c}                      & LMC  & ~9  & 280  &   8.37           & 30~Dor\tabnotemark{f}\\  
WR105\tabnotemark{a}                      & MW   & ~9  & 700  &                  &         \\
AG Car (during                            &      &     &      &                  &          \\
minimum epoch 1985--1990)\tabnotemark{d}  & MW   & 11  & 300  &                  &       \\
AG Car (during                            &      &     &      &                  &           \\
minimum epoch 2002)\tabnotemark{d}        & MW   & 11  & 200  &                  &       \\
\bottomrule
\tabnotetext{a}{\citet{smith}. \textsuperscript{b}\citet{b517}. \textsuperscript{c}\citet{Crowther1995}. \textsuperscript{d}\citet{AGcar}. 
\textsuperscript{e}\citet{OHabundance}. \textsuperscript{f}\citet{OHabundanceDor}.}
\end{tabular}
\vspace{0.3cm}
\end{table*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


\subsection{Nebular Lines}
\label{sec:nebula}

The lines at 4959 and 5007~\AA\ are clearly identified as [O~III] emissions, which is consistent with their flux ratio equal to 3 and FWHMs 
equal within the measurement errors. In the cooler spectra of 2003 \citep{polcaro}, nebular lines of [O~III] are overridden by N~II~4994-5005 
emissions. However, in 2007, the N~II lines are absent, but the [O~III]~$\lambda~4959,~5007$ doublet is clearly seen.

Nebular lines of [O~III]~$\lambda\lambda~4959,~5007$, [N~II]~$\lambda\lambda~6548,6583$ are present in all the spectra analysed by us while the 
[S~II]~$\lambda~6717,6731$  doublet is not. Therefore, we can estimate the electron density of the surrounding nebula. The electron density is 
between $10^{3.6}$~cm$^{-3}$ and $10^{4.9}$~cm$^{-3}$. The former value corresponds to the critical density of the [S~II] doublet, the second to 
that of [N~II]~$\lambda\lambda~6548,6583$ lines (see for example Osterbrock \& Ferland 2006).

All the nebular lines present in the FOCAS spectrum as well as the [Ar~III]~$\lambda~7135.73$ emission in SCORPIO data show flat-topped or 
two-peaked profiles. The widths of the emission line cores are of the order of 100~km~s$^{-1}$. Once we know the density of the emitting gas, we may 
estimate the size of the Str\"omgren region and the mass of the nebula surrounding V532. The former may be found as follows (see for example 
Lang 1974): 
$$R_\mathrm{s}\simeq \left(\frac{3S}{4\pi \alpha n_\mathrm{e}^2}\right)^{1/3}\,,$$
where $\alpha$ is the Case B recombination coefficient for hydrogen, $\alpha \simeq 2.6 \cdot 10^{-13} \left(\frac{10^4}{T}\right)^{0.85}$~cm$^3$~s$^{-1}$, 
$n_\mathrm{e}$ is the electron density, and $S$ is the hydrogen-ionizing quanta production rate. For O9.5Ia stars, probably similar to the object 
in mass and luminosity, $S$ is estimated as $10^{49.17}$~s$^{-1}$ (Osterbrock \& Ferland 2006).

One may estimate the size of a homogeneous nebula around V532 as:
\begin{eqnarray}
\label{E:rstrom}
R_\mathrm{s}\simeq 0.1  \cdot \left(\frac{S}{10^{49}{\rm s^{-1}}}\right)^{1/3} \times \left(\frac{n_\mathrm{e}}{10^{4}~\mathrm{cm}^{-3}}\right)^{-2/3} \times \nonumber \\
\times \left(\frac{T}{10^4\rm K}\right)^{0.28}~\mathrm{pc}\,.~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~
\end{eqnarray}

Note that possible inhomogeneity has little effect on the linear size: for a gas filling only some part of the nebular volume $f$, $R_\mathrm{s}$ should 
scale with filling factor as $\propto f^{-1/3}$. However, sizes of about 1~pc are plausible. The nebular mass may be trivially inferred as: 
$$M \simeq \rho \cdot \frac{4 \pi}{3} R_\mathrm{s}^3~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~$$
\begin{equation}
\label{E:mstrom}
\simeq 0.17 \cdot \left(\frac{S}{10^{49}{\rm s^{-1}}}\right) \left(\frac{n_\mathrm{e}}{10^{4}~\mathrm{cm}^3}\right)^{-1} \left(\frac{T}{10^4\rm K}\right)^{0.84} \cdot M_{\odot}\,. 
\end{equation}

\vspace{0.15cm}

The estimated size and the ionised mass of the nebula are consistent with these for ejecta of LBV stars (see Smith et al. 1994, and references therein) 
to order of magnitude. It is possible that the total mass lost by the object is higher, about several solar masses, but we observe only the ionized 
part of the envelope. The emitting gas was probably ejected during one or several outburst events at wind velocities of about 100~km~s$^{-1}$. 


%%%%% TABLA 6 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\begin{table*}[!t]
\vspace{0.1cm}
\centering
\caption{Equivalent widths of helium lines (in H$\beta$ units)\\ and helium-to-hydrogen abundance ratios 
for two\\ late hydrogen-rich WNs and V532 (2007)}
\label{tab:ew}
\begin{tabular}{llccccccc}
\toprule
Object         & Spectral & H/He  &  &  He~I+H8 & H$\gamma$  &  He~II  &   He~I    &  He~I  \\
               & Class    &       &  &  3889    &            &  4686   &   4713    &  5876 \\
\midrule
R84            & WN9      & 2.5   &  &   3~~    & 2.3~       &  3.5    &   12.6    &       \\
BE381          & WN9h     & 2~~   &  &   2.1    & 2.4~       &  1.9    &   16.6    &  0.8  \\
V532           &          &       &  &   2~~    & 2.75       &  1.4    &   13.4    &  0.8  \\
\bottomrule
\end{tabular}
\vspace{0.3cm}
\end{table*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%             


\vspace{0.15cm}

\section{Discussion}
\label{sec:disc}

In Table~\ref{tab:windvelocity}, terminal wind velocities for the WN9-WN11 stars in the Galaxy, LMC and M33 and LBV star AG Car (during minimum) 
are given. These velocities are also estimated through the optical He~I lines. Because wind acceleration is tightly connected to metallicity (Massey 
et al. 2004; Puls, Vink, \& Najarro 2008), we also give oxygen abundances (12 + log O/H) for selected H~II regions adjacent to MCA1-B, B517 and BE381. 
They may be used as a measure for initial object metallicities. 

V532 is located at a distance of about $17'$ from the centre of M33. Taking the central oxygen abundance value of $12+\log{\rm O/H}=8.36\pm0.04$ and 
a radial gradient of $-0.027\pm0.012$~dex~kpc$^{-1}$ (Rosolowsky \& Simon 2008), we may estimate the primordial oxygen abundance for V532 as 
$8.26\pm0.08$. The ambient metallicity is thus similar to that of the 30~Dor star-forming region. Our estimate for the terminal wind velocity of V532 
is fairly consistent with that of other late WN stars in similar environments.

Wolf-Rayet stars are believed to be more evolved objects than Ofpe and Be supergiants. Information on elementary abundances may cast some light 
upon this issue. However, to recover reliable abundance estimates, one needs sophisticated modelling of both the moving atmosphere of the star 
and its structure. Here, we restrict ourselves only to some semi-qualitative estimates for the H/He abundance ratio. We compare the equivalent 
width ratios of H$\beta$ to helium lines in the spectra obtained in 2007 to those for several WN9 (taken from Crowther et al. 1995) to estimate 
the relative abundances of hydrogen and helium. The equivalent width ratios are given in Table~\ref{tab:ew}. The table shows that they are similar 
to those for BE381 (Brey 64) that has H/He=2  and is classified as WN9h by \citet{CrowtherWNh}. By analogy with BE381, we suppose that for V532, 
H/He $\sim 2$. It places Romano's star in the region occupied by hydrogen-rich late WN stars that have atmospheres moderately contaminated by helium.

We have already mentioned the strikingly diverse behaviour of triplet lines of He~I having P~Cyg-like profiles and singlet emissions. Probably 
singlet lines have weak absorptions that we cannot detect due to insufficient signal-to-noise or spectral resolution. Besides this, He~I~$\lambda~5015$ 
is contaminated by the oxygen [O~III]~$\lambda~5007$ emission.

The similarity of the line profiles of singlet He~I, [O~III]~$\lambda\lambda~4959,~5007$ and [Ar~III]~$\lambda~7136$ suggests that both singlet He~I and 
forbidden emissions are produced mainly in the low-density ejecta, expanding at velocities of about 100~km~s$^{-1}$. Probably, singlet and triplet 
lines of neutral helium are formed in different places. Because the lowest possible level for triplet lines is metastable, concentrations of atoms 
at the lower levels of triplet transitions (and therefore the intensities of absorption line components) should be higher. 

We classify the observed evolution of the object as an S~Dor variability cycle. The main difference from AG~Car and other LBVs is that at the optical 
minimum V532 becomes much hotter and may be classified as a WN9/WN8 star, while AG~Car stops at about WN11. The non-monotonous spectral changes and 
variability timescales indicate that these spectral changes hardly correspond to any real stellar evolution, rather being connected to the ordinary 
LBV variability. It also means that LBV stars may have spectral classes as early as WN8. To our notion, there is only one example of an object that 
acquired even earlier spectral classes of WN6/7 during LBV variability cycle, HD5980 (Barb\'a, Niemela, \& Morrell 1997; Koenigsberger \& Moreno 2008) 
in the SMC. The hot spectra of this object are probably connected to lower ambient metallicity, too. 

V532 is one more example of an LBV star temporarily becoming a WN. We predict that, vice versa, some of the Galactic or extra-galactic Wolf-Rayet 
stars may prove to be dormant LBVs. Probably, the hottest possible temperature for an LBV decreases with metallicity, but more data on the brightest 
stars in different environments are needed. 

\vspace{0.2cm}

\section{Conclusions}
\label{sec:con}
     
Our results show that the object changes from a $B$ emission line supergiant at optical maximum, through Ofpe/WN (WN10,WN11) to WN9 and further 
towards a WN8 star at deep minimum. We confirm the result of \citet{viotti2007} that there is an anti-correlation between the visual luminosity 
and the temperature of the star, with a larger amount of spectral data. 

V532 spans a large range of spectral classes, becoming one of the first (probably, the second, after HD5980) LBV stars noticed to make an excursion 
as deep as WN8 into the Wolf-Rayet domain. It is interesting to trace further the evolution of the object. Further monitoring of V532, as well as 
of other late WN stars, is necessary to establish the evolutionary link between WRs and LBVs. Some WNs may prove to be dormant LBVs. 

\vspace{0.15cm}

\acknowledgements

We are grateful to Vitalij Goranskij, Elena Barsukova and Alla Zharova for providing us with photometric data and Olga Sholukhova and 
Thomas Szeifert for the spectrum obtained with TWIN Calar Alto spectrograph in 1992. We would also like to thank the anonymous referee 
for valuable comments and for drawing our attention to HD5980. 

%\vspace{0.2cm}

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\end{document}



