% 1999-04-29: Manuscript completed for the AJ and submitted.
% 
% 1999-05-25: Revised version; includes the changes reqested by the referee.

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\lefthead{Tornikoski et al.}
\righthead{Multiwavelength observations of PKS~2255$-$282}

\begin{document}

\def\K{\,{\mathrm{K}}}
\def\micron{\,{\mu}{\mathrm{m}}}
\def\mas{\,{\mathrm{mas}}}
\def\mag{\,{\mathrm{mag}}}

\title{Multiwavelength observations of PKS~2255$-$282}

\author{M. Tornikoski, \altaffilmark{1} % Merja.Tornikoski@Hut.Fi
        S.J. Tingay,   \altaffilmark{2} % tingay@hyaa.jpl.nasa.gov
        A. M\"ucke,    \altaffilmark{3} % amuecke@physics.adelaide.edu.au
        A. Chen,       \altaffilmark{4} % awc@blazar.gsfc.nasa.gov
        V. Connaughton,\altaffilmark{9,11} % vc@msfc.nasa.gov
        D.L. Jauncey,  \altaffilmark{5} % djauncey@atnf.csiro.au
       M. Johnston-Hollitt, \altaffilmark{3} % mjohnsto@physics.adelaide.edu.au
        J. Kemp,    \altaffilmark{6,12} % jonathan@cba.phys.columbia.edu
        E.A. King,      \altaffilmark{5}   % eking@atnf.csiro.au
        P. McGee,       \altaffilmark{3}   % pmcgee@physics.adelaide.edu.au
        F. Rantakyr\"o, \altaffilmark{7}   % frantaky@eso.org
        D. Rayner,      \altaffilmark{8}   % drayner@kerr.phys.utas.edu.au
        O. Reimer      \altaffilmark{10}  % olr@mpe-garching.mpg.de
   and  A.K. Tzioumis  \altaffilmark{5}   % atzioumi@atnf.csiro.au
}

\altaffiltext{1}{
           Mets\"ahovi Radio Observatory, Mets\"ahovintie 114, 
           FIN--02540 Kylm\"al\"a, Finland;
           E-mail: Merja.Tornikoski@Hut.Fi}
\altaffiltext{2}{
           Jet Propulsion Laboratory (MS238-332), 
           California Institute of Technology, 
           4800 Oak Grove Drive, Pasadena, CA 91109, USA}
\altaffiltext{3}{
           University of Adelaide, 
           Dept. of Physics and Math. Physics,
           Adelaide, SA 5005, Australia}
\altaffiltext{4}{
           NASA/Goddard Space Flight Center, Greenbelt, MD 20771, USA}
\altaffiltext{5}{
           Australia Telescope National Facility,
           P.O. Box 76, Epping, NSW 2121, Australia}
\altaffiltext{6}{
           Biosphere 2 Center, Columbia University, 32540 S Biosphere Rd,
           Oracle, AZ 85623, USA}
\altaffiltext{7}{
           European Southern Observatory / SEST, 
           Casilla 19001, Santiago 19, Chile}
\altaffiltext{8}{
           University of Tasmania, Dept. of Physics,
           G.P.O. Box 252C, Hobart, Tasmania 7001, Australia}
\altaffiltext{9}{
           NASA Marshall Space Flight Center, AL 35812, USA}
\altaffiltext{10}{
           Max-Planck-Institut f\"ur extraterrestrische Physik,
           Postfach 1603, D-85740 Garching, Germany}
\altaffiltext{11}{
           NAS/NRC Research Associate}
\altaffiltext{12}{
           Visiting Astronomer, Cerro Tololo Inter-American Observatory,
           National Optical Astronomy Observatories, which is operated by the
           Association of Universities for Research in Astronomy, Inc. (AURA)
           under cooperative agreement with the National Science Foundation.}

%-------------- ABSTRACT:

\begin{abstract}

During high radio frequency AGN monitoring we observed a major
millimeter-wave outburst from PKS~2255$-$282 in 1997 and initiated a
campaign for multifrequency radio monitoring and VLBI imaging of the
source.  In January 1998 the EGRET instrument onboard the CGRO
detected a bright high-energy $\gamma$--ray outburst from its
direction. Thus we had the unique opportunity to monitor the source in
detail at various radio frequencies already before the detection of
the $\gamma$--ray outburst. We also initiated multiwavelength followup
observations. In this paper we present multiwavelength data for
PKS~2255$-$282 before and after the $\gamma$--ray outburst.  Further,
based on the observations, we estimate some physical parameters for
this source and discuss the various emission mechanisms possibly
involved during this outburst.
\end{abstract}

\keywords{quasars: general --
          quasars: individual (PKS~2255$-$282)}


%-------------- TEXT:

\section{Introduction}

The radio-loud active galactic nucleus (AGN) PKS~2255$-$282, z=0.926,
is a southern source that did not receive very much attention prior to
its flaring at $\gamma$--ray energies in January 1998.  Historical
observations of PKS~2255$-$282
(Bolton \& Butler \cite{bolton}; Condon et al. \cite{condon}; 
K\"uhr et al. \cite{kuehr}; Large et al. \cite{large}; 
Morabito et al. \cite{morabito}; the Parkes catalog \cite{pks}; 
Robertson et al. \cite{robertson}; Quiniento \& Cersosimo \cite{qui:cer};
Steppe et al. \cite{steppe93}; Griffith et al. \cite{griffith},
and the historical SEST--radio measurements (Tornikoski et al.
\cite{tornikoski96} and this paper)) show that it is a
bright flat-spectrum radio source, with no prominent radio or optical
outbursts reported in the literature.  
Optical polarization measurements conducted in June 1986 (Impey \&
Tapia \cite{imp:tap}) gave a polarization of $2\% (\pm0.4\%)$ and
polarization angle of $112\arcdeg (\pm6\arcdeg)$. 
This classifies PKS~2255$-$282 as a low polarization quasar (LPQ), based 
on the criterion where LPQs have a degree of polarization
$p<3\%$ and for high polarization quasars (HPQ) $p>3\%$.

Millimeter-wave radio data for PKS~2255$-$282 have
been collected since the end of 1987 as part of the SEST AGN
monitoring project undertaken by the Mets\"{a}hovi group (Tornikoski
et al. \cite{tornikoski96}). Classified as one of the lower priority
sources, the data set for this source is initially relatively sparsely
sampled. No major outbursts were observed before 1997.

When the mm--flare was observed in 1997, a campaign for more frequent
monitoring of the source at several radio wavelengths was initiated.
During this campaign, EGRET observations onboard the Compton Gamma-Ray
Observatory (CGRO) detected a very bright $\gamma$--ray flare in
PKS~2255$-$282.  

In the sections 2 to 6 we present multifrequency data obtained before,
during and after the $\gamma$--ray flare.  In section 7 we estimate
various physical parameters from our data, (including several
estimates of the Doppler factor of the emitting material) and discuss
the multifrequency variability behaviour within the framework of
common high-energy emission models.

\section{Radio variability}

\subsection{Observations}

At millimeter wavelengths PKS~2255$-$282 was observed as part of the
AGN monitoring project with the 15-meter Swedish-ESO Submillimeter
Telescope (SEST) at the European Southern Observatory site on Cerro La
Silla in Chile.  The Mets\"{a}hovi group has monitored the flux
variability of a relatively large sample of AGN at millimeter
wavelengths since the end of 1987.  The observations have been made at
3 mm / 90 GHz and 1.3 mm / 230 GHz.
Until June 1995 the 3 mm observations were made
using a dual polarisation Schottky receiver, and after that the
measurements were made with an SIS receiver.  
As a backend, a wide-band (1 GHz) acousto-optic spectrometer was
used. For the 1.3 mm observations we also initially used a Schottky
reciever and a wide band AOS, but since 1991 measurements have been
mainly obtained with a single channel bolometer. All the flux density
measurements were made in a dual beamswitching mode, and the flux
densities were calibrated against planets. 
Details about the SEST monitoring project as well as the full data
base of the project until June 1994 can be found in Tornikoski et
al. (\cite{tornikoski96}).

At centimeter wavelengths PKS~2255$-$282 was observed as part of an
on-going monitoring program undertaken with the Australia Telescope
Compact Array (ATCA). Observations were made at 1.6, 2.4, 4.8 and 8.6
GHz. Data were recorded in XX and YY polarisations using bandwidths of
104 MHz and scan lengths of 1--2 minutes per observing frequency.  The
data were calibrated and reduced using the MIRIAD software package
(Sault \& Killeen \cite {sau:kil}). To estimate the flux density of
PKS~2255$-$282, a point source was fitted to the data. The error bars
on the flux density are therefore dominated by the extended structure
of the source (the effects of which are small since the source is
highly compact). The errors in the flux density scale were estimated
to typically be between 1 and 2\%.  The data points from 1998 January 19
and 27 were based on 15 minute and 5 minute scans at 4.8 GHz
and 8.6 GHz, respectively, with a bandwidth of 128 MHz. The error bars
for these data points were 5--6\% of the flux density.  PKS~1934-638
served as a primary calibrator for these data points.

The centimeter wavelength data points presented in this paper
were collected between July 1996 and August 1998.

\subsection{Radio variability from 1987 to 1997 (before the mm--flare)}

Our SEST monitoring data for PKS~2255$-$282 before 1997 is relatively
sparse (between 1987 and 1996 only 1--2 SEST 90 GHz data points per
year), but we have complemented our data sets with IRAM data from
the literature (Steppe et al. \cite{steppe88}, \cite{steppe92},
\cite{steppe93}; Reuter et al. \cite{reuter}). The average radio flux 
densities from this period are given in Table 1.
The 90 and 230 GHz flux curves with SEST and IRAM data are shown in 
Fig.~\ref{90_230_EGRET}.


In spite of the somewhat sparse sampling we can see no indication of
any very rapid and strong outbursts before the 1997 flare. In
1988--1989 there was a moderate flare observed, during which the
90 GHz flux increased from ca. 2.5 Jy to the maximum of 5 Jy,
and after 1990 the flux level remained at a relatively constant value
slightly below 2 Jy. 

Towards the end of 1994 there were already indications of the mm--flux 
starting to rise, but flux levels ca. twice the quiescent flux values
(4 Jy at 90 GHz) were not reached until 1996.

We unfortunately had no lower radio frequency monitoring for
PKS~2255$-$282 during the moderate flare in 1988--1989 or the
quiescent period in 1990--1996, but we have collected data from
the literature (see the references listed in Section 1)
indicating that no major flares have been observed in the cm--wave
domain.


\subsection{Radio variability from 1997 to 1998 (during the mm--flaring 
            period)}

In April 1997 the flux density of PKS~2255$-$282 at 90 GHz
was observed to be over 8 Jy, twice the flux density observed 11 months
earlier. This was a clear indication of a very strong mm--wave
outburst and alerted our group for multiwavelength radio
monitoring and more dense sampling.

At the end of June 1997 the 90 GHz flux was 10.1 Jy, making
PKS~2255$-$282 one of the very small handful of sources to ever have
exceeded 10 Jy at 3 mm. By August 1997 the flux density was
back to $55\%$ of the peak value at 230 GHz and $65\%$ at
90 GHz, but there was a moderate rise in the 90 GHz flux again towards
the end of the year.

The multifrequency radio flux curves since the beginning of 1996 are
shown in Fig.~\ref{flareplot_radio}. At 8.6 GHz we observed a notable
increase in the flux level simultaneously (within our sampling 
accuracy) with the higher frequency flux. A moderate
increase from the 1996.5 flux levels were also observed at 4.8 GHz and
2.3 GHz, while the 1.6 GHz data showed almost constant flux densities
throughout the outburst.

After the peak of the mm--flare was observed, also the 8.6 GHz flux
density decreased but from the beginning of 1998 there has been
another period of flux increase which still continued in August.  
By August 5th, 1998, the 8.6 and 4.8
GHz flux curves had reached their all time maxima: 6.6 Jy at 8.6 GHz
and 3.6 Jy at 4.8 GHz. The new rising trend in the centimeter flux
curves coincides with the very bright $\gamma$--ray outburst observed
with the EGRET instrument in January 1998 from the direction of
PKS~2255$-$282. The radio flux densities during the
$\gamma$--ray outburst are given in Table 2.

In the mm--data there is a longer gap in 1998, mainly due to the SEST
Sun limit. In August 1998 the 90 GHz flux was observed to be at
approximately the same level as in December 1997.  
The behaviour of the 8.6 GHz flux curve as well as the still relatively
high 90 GHz flux levels in late 1998 can be seen as an indication of a
continuation of a very active period in the radio domain.
The radio flux densities from the followup obserations after the 
$\gamma$--ray outburst are given in Table 3.

When following the approach by Ter\"asranta and Valtaoja
(\cite{ter:val}) to calculate the source brightness temperature,
$T_{b,obs}$, and then the Doppler boosting factor from the radio
variability observations, we get a $T_{b,obs} < 10^{12}\K$ for the 90
GHz flare in 1997. This contradicts the $T_b$ estimate from the VLBI
data (Section 3) as well as the Doppler boosting factor estimates
obtained from the high-energy data (Section 7).  This can possibly be
explained by the severe undersampling of the mm--data: we can not
conclusively say when the flare starts or when it peaks. Another
possibility is that there are two superposed flare components in the
observed radio flux: one slowly increasing and decaying component, and
another very rapid one. Indications of such superposed events also in
higher frequency data have been found by Valtaoja et
al. (\cite{val:lah}) in the well-sampled Mets\"ahovi AGN monitoring
data. VLBI data obtained at later post-flare epochs may give us an
indication of the true number of the underlying shock components.

\subsection{Radio spectrum}

The radio spectrum of PKS~2255$-$282 is very flat, especially 
during the mm--outburst. The quiescent spectra (historical data and
epoch 1996.5 data) peak above 10 GHz, and the mm--flare spectrum
peaks around 90 GHz. The radio spectral indices $\alpha$ 
(S$\propto\nu^{\alpha})$ for the quiescent state are
$\alpha=-0.03$ betweeen 8.6 and 90 GHz and  $\alpha=-0.57$ between 
90 and 230 GHz,
and during the mm--flare they are  $\alpha=0.22$ and  $\alpha=-0.41$, 
respectively.
In a study by Tornikoski et al. (\cite{tornikoski93}) the average
spectral index between 8.4 and 90 GHz for southern sources during
their quiescent state was found to be  $\alpha=-0.25$ for HPQs (14 sources)
and  $\alpha=-0.46$ for LPQs (24 sources). We can see that even in the
quiescent state PKS~2255$-$282 has a radio spectrum flatter than
that of ordinary quasars and even many HPQs, and during the outbursts 
the spectrum still flattens, so the radio behaviour of this source
resembles that of the blazars much more than ordinary quasars.

\section{VLBI Observations}

VLBA observations of PKS~2255$-$282 were made on 1997 August 13 in the
5, 22, and 43 GHz bands.  The data were reduced in AIPS using the
standard set of reduction tasks for continuum VLBI and then imported
into the DIFMAP software (\cite{shepherd}) for model-fitting and imaging.
Fig.~\ref{VLBI_image} shows the resultant 5 GHz image, revealing a bright,
marginally resolved core component of 2.8 Jy and a short ($\sim5\mas$) 
jet extending to the south-west which contributes
approximately 0.25 Jy to the image.  At both 22 GHz and 43 GHz the jet
feature was completely resolved and absent from images, however the
core component at these frequencies contained 4.2 Jy and 2.7 Jy,
respectively.  Previously Fey, Clegg \& Fomalont  
(\cite{fey}; hereafter FCF) had
imaged PKS~2255$-$282, in 1994, at 2.3 and 8.4 GHz.  The structure
seen in Fig.~\ref{VLBI_image} is consistent with that of FCF at 8.4 GHz, but at
2.3 GHz FCF show that the jet feature extends approximately 20 mas
from the core.

Table 4 lists the models fit using the MODELFIT command in
DIFMAP, for our new 5, 22, and 43 GHz VLBA data, in the u-v plane.

Brightness temperature estimates for the core are also listed in 
Table 4, for the core component, in units of $10^{12}\K$, 
for Gaussian,
optically thin sphere, and optically thick sphere model components,
calculated in the source frame.  Even using the least centrally
concentrated component, the optically thick sphere, the brightness
temperature exceeds $10^{12}\K$ at 22 GHz, indicating that beaming of
the radio emission is required.

Moellenbrock et al. (\cite{moellenbrock}) measured the brightness
temperature from the observations using single baseline VLBI
observations at 22 GHz.  Using a Gaussian model component they
obtained a lower limit to the brightness temperature of $0.2 \times
10^{12}\K$.  The 1994 VLBA observations by FCF yield a brightness
temperature of $2.7 \times 10^{12}\K$ at 8.4 GHz, using a Gaussian
component.  For the optically thin and optically thick sphere cases,
this corresponds to brightness temperatures of $1.2 \times 10^{12}\K$
and $1.0 \times 10^{12}\K$, respectively.

From the observations of FCF and our new VLBA
observations at higher frequency, it appears that the nominal
brightness temperature limit of $10^{12}\K$ is being reached and
perhaps exceeded in PKS~2255$-$282.  Thus, beaming of the emission is
required, which is reasonable for a core dominated radio source such as
this.  For PKS~2255$-$282 to produce its high-energy emission in
$\gamma$--rays it would seem plausible that its intrinsic brightness
temperature is close to $10^{12}\K$, if inverse Compton
processes are responsible for the high-energy emission.  If this is the
case then the amount of Doppler boosting can be roughly estimated by
comparing a plausible intrinsic brightness temperature to the measured
brightness temperature.  An intrinsic brightness temperature close to
$10^{12}\K$ and a measured brightness temperature only a factor of 2 to
3 higher means that the Doppler factor, $\delta$, is likely to be 2--3.

With future VLBA observations at 5 GHz, we plan to measure any apparent
motion in the jet components of PKS~2255$-$282 and obtain further
constraints on the beaming properties of this source.


\section{Optical observations}

Historical optical observations for PKS~2255$-$282 have been 
published in Wright et al. (\cite{wright}), Impey \& Tapia
(\cite{imp:tap}) and Bowen et al. (\cite{bowen}).

We made optical follow-up observations for PKS~2255$-$282 in the
post-flare phase from July to September 1998 at the 0.35m telescope at
Woomera (equipped with a SBIG-ST6 CCD camera) and the 0.9m telescope
at CTIO.  The Woomera telescope used a standard R-filter, while
measurements at the CTIO-0.9m telescope were performed in standard VRI
bands.  A sequence of nearby reference stars (see Table 5
and Fig.~\ref{field}) in the field of PKS~2255$-$282 was
established at the CTIO telescope using standard stars from the list
of Landolt (\cite{landolt}), and were used for calibrating the
measurements of PKS~2255$-$282.  At the Woomera telescope,
differential photometric measurements of PKS~2255$-$282 were performed
with respect to these secondary standards. Multiple exposure (8 frames
with 4 min integration time each) were obtained to minimize the
observational uncertainties.  The results are given in
Table 6.

\section{High-energy observations}

During the CGRO observations from 1997 December 30 to 1998 January 12
(VP708/709), PKS~2255$-$282 was $\sim13\arcdeg$ from the center of the
field of view of EGRET and COMPTEL.  In the EGRET observation of
January 1998 a strong outburst, comparable with the 3C~279 flare from
June 1991, was detected, position-compatible with PKS~2255$-$282
(Bertsch et al. \cite{bertsch}).  Standard analysis of the summed
EGRET maps of viewing periods VP708/709 yields a $9.6\sigma$-detection
of PKS~2255$-$282. Its spectrum follows a straight power law between
30 MeV and 10 GeV with photon spectral index
$\alpha_{\gamma}=1.58\pm0.17$ (reduced $\chi^2=1.018$) and no
indication of a break at high energies. We use the convention $N_{ph}
\sim E^{-\alpha_{\gamma}}$ for the photon spectrum $N_{ph}$ and use
the measured spectral index to convert count rates in energy flux.
The integral photon flux $>$100 MeV reads $(57.1\pm10.0) \times
10^{-7} $cm$^{-2}$ s$^{-1}$.  The inspection of the arriving events on
a day-by-day basis reveals that the maximum energy output occured
around day 9/10 of this viewing period though the most energetic event
($\simeq 11$ GeV) arrived on day 2.  By the end of this viewing period
the count rate had dropped by about a factor of 3.
The variability time scale lies in the order of
several days.  Historically, PKS~2255$-$282 was below the detection
threshold of EGRET from 1991--92, yielding a $1\sigma$ upper limit of
$1.48 \times 10^{-7}$cm$^{-2}$ s$^{-1}$ ($>$100 MeV). In November 1997
EGRET measured a flux of $(4.7 \pm 2.3)\times10^{-7}$cm$^{-2}$ s$^{-1}$
($>$100 MeV), a $3\sigma$-detection, and roughly 10 times smaller than
the flux from the January 1998 outburst.
The differential EGRET gamma-ray energy spectrum during the January
observations is given in Table 7.

The analysis of the BATSE data in the energy range 20--200 keV revealed no
significant signal from PKS~2255$-$282 during the time of the EGRET-outburst.
The $2\sigma$ upper limit for the time integrated
signal (from 1997 December 30 till 1998 January 12) is calculated by
folding the time integrated photon flux through the response function
with a $-2$ power law photon spectral index, and is included in 
Fig.~\ref{SED} .

COMPTEL has not seen this source, both during its EGRET-flaring state
(Collmar \cite{collmar}) and in previous observations. Therefore, we
use the $3\sigma$ sensitivity limits for a typical two-week
observation run (Kurfess et al. \cite{kurfess}) to constrain the SED
in the COMPTEL energy range (0.75 MeV -- 30 MeV) (see
Fig.~\ref{SED}).


The historical ROSAT measurement point, taken during the ROSAT allsky
survey (from 1990--91), gives $(9.43\pm3.25)\times 10^{-13}$ erg
cm$^{-2}$ s$^{-1}$ in the 0.1--2.4 keV band, and is published in
Brinkmann et al. (\cite{brinkmann}).  
Probably due to the high HI column density in the
direction of PKS~2255$-$282, this source was not detected either by
the IUE or the EUVE-satellite.

XTE observations of PKS~2255$-$282 were conducted soon after the
January 1998 $\gamma$-ray flare, on January 15th, and yield a flux
level of $(9.0\pm1.4)\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ in the
2--10 keV range. This is roughly an order of magnitude higher than the
historical ROSAT measurement point.

Data reduction and analysis of the high-energy observations will
be given in Chen et al. (\cite{chen}).  


\section{The spectral energy distribution (SED)}

PKS~2255$-$282 is one of the few EGRET blazars with a fairly good
coverage at cm-- and mm--radio wavelengths prior to its first
detection at $\gamma$--ray energies by EGRET. Therefore we were able to
construct its spectral energy distribution in different states of
activity.  Since the mm--radio data base of PKS~2255$-$282 is by far
the most complete one compared to the other wavelengths, we chose this
band to define the different temporal phases except for the flaring
state. 

We defined the following four activity states for our SED study:
\begin{enumerate}
\item The {\em quiescent state}, defined here as the average intensity between
      January 1990 and December 1993 in addition to all historical data
      prior to 1986.
\item The {\em pre-flare state} January 1994 to December 1997 before the
      $\gamma$-ray flare was observed.
\item The {\em flare-state}, defined as the time span of the CGRO observation 
      during the $\gamma$--ray flare, plus one week.
\item The {\em post-flare state}, data collected in 1998 after the CGRO
      observations were ended.
\end{enumerate}

Fig.~\ref{SED} shows the energy distribution $\nu f_{\nu}$ for the
quiescent, pre-flare, flare and post-flare state.

Only radio measurements are available for the pre-flare state.  The
flux densities show a continuous increase by a factor 4--5 over this
state.  The average pre-flare flux densities are given in Table 1. The
flare-state includes cm--, mm-- XTE-- and EGRET--measurements (see
Table 2 and Table 7).

In the radio regime the SED (Fig.~\ref{SED}) is well covered and shows a
hardening with increasing flux density. This can be understood as the
result of an increasing variability amplitude with increasing
frequency. We note that the 4.8 GHz and 8.6 GHz flux densities in the
post-flare state are still slightly rising, indicating a variability
time scale of years.  Unfortuntely, no data in the far-infrared
frequency band are available, which prevents us from determining
conclusively the maximum of the synchrotron hump in the SED. According
to Fossati et al. (\cite{fossati}), the maximum synchrotron power in FSRQs is
expected to be emitted at $\sim10^{12}-10^{14}$ Hz, which is not in
conflict with our observations.

Optical photometry was only obtained in the post-flare state.
Compared to the historical measurements 
only a slight brightening of $\sim 0.5$mag is noted. 
This would mean either that large-amplitude variability in 
the optical band happens on time scales less than months, 
or that only small-amplitude variations occur. 
The apparent positive slope of the historical NIR/optical data
points in the $\nu f_{\nu}$-plot may be due to observations at
different activity states, while the simultaneous optical post-flare
state measurements might indicate the onset of a ``blue-bump''
observed in a number of blazars. However, this can not be stated
conclusively due to sparse frequency coverage in this spectral
range. Note also that strong broad emission lines are observed from
PKS~2255$-$282 (e.g., Wilkes \cite{wilkes}; Peterson et
al. \cite{peterson}).

In the high-energy region a second hump in the SED is expected. In
PKS~2255$-$282 the peak of this component must lie at energies $>$ 10
GeV as indicated by the hard spectrum in the EGRET-energy range. 
The similar power law spectrum in the XTE- and EGRET energy range
together with the fact that the XTE-data points lie readily on the
extension of the EGRET-spectrum within the observational uncertainties
suggest the same origin for the X--ray and $\gamma$--ray emission.
The dramatic increase of the GeV flux during the outburst, and the
degree to which the power in this waveband dominates the $\nu f_{\nu}$
spectrum, is consistent with the behaviour found in other
EGRET-detected flat spectrum radio sources (e.g. von Montigny et
al. \cite{montigny}).

\section{Discussion}

The observational multifrequency data base we have obtained for PKS 2255$-$282,
complemented with data from the literature, extends from the low radio
frequencies to the $\gamma$--ray energies, and our longest radio data
sets span more than 10 years. Even though our sampling has
been somewhat sparse and in most cases the multifrequency coverage has
been only semi-simultaneous, we have been able to construct
multifrequency spectra and spectral energy distribution diagrams for
several epochs and activity states.  Furthermore, this data set for
the first time includes a significant amount of monitoring data {\em
prior} to a major $\gamma$--ray flare.

The $\gamma$--ray flare observed in PKS~2255$-$282 in January 1998 was
observed during a period of increased mm--wave activity.
PKS~2255$-$282 has been in the EGRET field of view several times
before January 1998, as indicated in Fig.~\ref{90_230_EGRET}, but was
not detected in $\gamma$--rays until the source was in an active
mm--state. This clearly indicates that the $\gamma$--ray flare was
connected with the brightening in the mm--wave bands.  Though the peak
fluxes at 90 and 230 GHz were already reached 6 months before the
$\gamma$--ray flare was observed, the multifrequency radio flux curves
indicate that in January 1998 the source was still in a very active
state also in the radio domain.  So it appears that radio activity in
PKS~2255$-$282 is related to $\gamma$--ray emission in the EGRET-energy
range.  Furthermore, the $\gamma$--ray outburst occured during the
increasing phase of the cm--radio emission.  Valtaoja and Ter\"asranta
(\cite{val:ter}) have argued that the detection of $\gamma$--ray flares
during the onset of high radio frequency flares is a clear indication
that $\gamma$--rays originate from the same shock as the synchrotron
radio emission. The connection between activity in the radio domain
and the EGRET energy domain for active EGRET--blazars was also found
by M\"ucke et al. (\cite{muecke}) on a statistical basis.  However,
due to the undersampling of the $\gamma$--ray data for this source, a
direct luminosity correlation between radio and $\gamma$--ray emission
can not be confirmed.

Rapid time variability in the $\gamma$--ray regime together with
co-spatially emitted GeV and keV photons can place strong constraints
on the beaming factor of the high-energy radiation. For PKS~2255$-$282
the observed GeV and keV fluxes, which we assume to originate co-spatially,
imply that the optical depth due to photon-photon pair
production can not exceed unity for this source.  With this
constraint, and using a variability time scale of several days in the
$\gamma$-ray band we can estimate a lower limit for the Doppler factor
(see e.g. Mattox et al. \cite{mattox}, Eq.4; note, however, that in
this equation a factor $(1+z)^4$ is missing),
$$
\delta \approx 3 (\frac{E_{\gamma}}{\rm{GeV}})^{0.118} ,
$$
where $E_{\gamma}$ is the observed $\gamma$--ray energy. For
$E_{\gamma} = 0.03-10$ GeV the minimum Doppler factor is $\delta =
2-4$, in approximate agreement with the ``inverse Compton
catastrophe'' limit derived from the VLBI observations.  Note, that
neither of these two limits for the Doppler factors depend on a
particular model for the high-energy emission.

Simultaneous multifrequency observations may be able to discriminate
between the different models for the high-energy emission.  From the
radio to the UV--band (or even X--ray band in some TeV-blazars)
relativistic electrons radiate synchrotron photons, and at higher
energies low energy photons are up-scattered to $\gamma$--ray energies
by the same population of electrons.
In the synchrotron self-Compton (SSC) model the target photons for the
inverse Compton (IC) scattering originate from the synchrotron photons
generated in the jet, and the IC component of the spectrum has the same
spectral shape as the synchrotron component. In the SSC model
one expects the optically thin synchrotron emission to vary 
simultaneously, but unequally in amplitude, with the high-energy radiation.  
In the external inverse Compton (EIC) model the
target photons for the IC scattering originate from a photon source
external to the jet: an accretion disk (e.g. Dermer \& Schlickeiser
\cite{dermer}) or reprocessed radiation from the surrounding BLR
and/or the dusty torus (e.g.  Sikora et al. \cite{sikora}, Ghisellini
\& Madau \cite{ghi:mad}).  If the $\gamma$--ray flare is caused by an
increase in the electron density in the emerging plasma blob, the EIC
model also predicts simultaneous variations of the synchrotron and IC
component, but in contrast to the SSC model with comparable amplitudes
in both high and low energy components.  Unfortunately, neither the
luminosities of the synchrotron and IC peak nor the exact epochs of
the radio and high-energy peaks can reliably be estimated from our
data for PKS~2255$-$282. 

In the EIC model variations can also be caused by changes in the seed
photon flux which should be observed before the $\gamma$--ray flare. In
this case, the synchrotron and IC component would not show correlated
variations.  However, IR to UV data are not available for
PKS~2255-282, preventing us from investigating this possibility.  We
conclude that, on the basis of our data, it is not possible to decide
which of the SSC or EIC mechanisms is dominantly responsible for the
high-energy emission from PKS~2255$-$282.

Adopting a simple one-zone homogeneous SSC model, we can derive an
estimate for the beaming factor if the observed X--ray emission is due
to the SSC mechanism (e.g., Jones et al. \cite{jones}; Marscher
\cite{marscher}).  We assume that the electron spectrum (in the jet
frame) is distributed isotropically in a spherical plasma blob and
follows a power law $\propto \gamma^{-2\alpha-1}$ where
$\gamma$ is the Lorentz factor of the electrons and $\alpha$ the
energy spectral index of the optically thin synchrotron spectrum ($F
\propto \nu^{-\alpha}$). Furthermore, the magnetic field is assumed to
be homogeneous and isotropic. Following Ghisellini et
al. (\cite{ghi:mad}) the minimum Doppler factor for the case of a
moving sphere is then given by
$$
\delta = f(\alpha)F_m (1+z) [\frac{\ln(\nu_b/\nu_m)}{F_x 
\Theta_d^{6+4\alpha}\nu_x^\alpha
\nu_m^{5+3\alpha}}]^{1/(4+2\alpha)}
$$
where: $f(\alpha) = 0.08\alpha+0.14$; $F_m$ is the radio flux from the
core; $\Theta_d$ is the angular size of the core at the
synchrotron-self absorption frequency $\nu_m$; $\nu_b$ is the
synchrotron high frequency cutoff; and $F_x$ is the X--ray flux at
frequency $\nu_x$. From our VLBI images we derive
$\Theta_d\approx0.28\mas$, $F_m=4.2$ Jy at $\nu_m=22$ GHz and
$\alpha=0.66$ the spectral index between 22 GHz and 43 GHz. We
estimate the synchrotron cutoff to be $\nu_b \sim 10^{13}$ Hz.  With
these numbers we derive $\delta\sim5$ and 7 using the 
XTE and ROSAT X--ray flux, respectively. This is roughly a factor of 2 higher
than our previous estimates for the beaming factor, however, it is
model dependent and also suffers from the non-simultaneity of the
radio and X--ray data.

In the framework of a homogeneous SSC model the low-energy component is
due to synchrotron radiation and the high-energy component due to
inverse Compton (IC) scattering of the low energy component by the
same population of electrons. The observed SED (Fig.~\ref{SED}) suggests
that the synchrotron peak lies around $\nu_s \sim10^{13}$ Hz and the IC peak
is at $\nu_{IC} >10^{24}$ Hz. The IC scattering producing the GeV
$\gamma$--rays mainly lies in the Thomson regime since $\nu_{IC} \leq
(\delta mc^2/h)^2/\nu_s$. The Lorentz factor of the radiating
electrons in the synchrotron/IC-hump can thus be estimated to be
$\gamma_b^{\prime} = \sqrt{\frac{3\nu_{IC}}{4\nu_{s}}} \sim 3\times
10^5$.  Here and in the following, primed quantities refer to the jet
frame.  The maximum energy of the synchrotron component is determined
by the magnetic field $B^{\prime}$ and the energy of the electrons
$\gamma_b^{\prime}$. Thus, we can estimate the magnetic field with
$B^{\prime} \simeq [(1+z)/\delta] [\nu_s^2/2.8\times 10^6 \nu_{IC}]
\leq (1-2) \times 10^{-5}$ G for $\delta = 2-7$.


\section{Conclusions}

We have collected multifrequency data for PKS~2255$-$282 before,
during and after the bright $\gamma$--ray outburst observed in 
January 1998. 

The radio spectrum of PKS~2255$-$282 is already relatively flat in 
the quiescent state, but it still flattens during the active state.  
The radio spectrum and variability resemble those of blazars.  Optical
observations were made only in the post-flare phase and show no
indications of optical flaring.

In January 1998 the source underwent a short period of activity
in the high-energy domain, detected in the EGRET $\gamma$--ray
and the XTE X--ray observations. We have carried out data analysis 
based on our quasi-simultaneous multifrequency data. Even though 
the activity in the radio domain appears to be related to the
$\gamma$--ray emission, which supports the SSC model, we can not
rule out the possibility of the EIC mechanism, either.
The similar power law spectrum in the XTE and EGRET energy range
suggest the same origin for the X--ray and $\gamma$--ray emission.
The peak of the high-energy component is expected to lie at energies
$>$ 10 GeV.

After the first version of this paper was finished, a study on the 
$\gamma$--ray flaring of PKS~2255$-$282 was published in Macomb,
Gehrels \& Shrader (\cite{macomb}). The EGRET data analysis of Macomb 
et al. differs slightly from the usual method, which was also used in
this paper, and yields a total $\gamma$--ray flux of $>$100 MeV  
$(1.6\pm0.3) \times 10^{-6} $cm$^{-2}$ s$^{-1}$ with a peak flux
of $(4.8\pm1.1) \times 10^{-6} $cm$^{-2}$ s$^{-1}$.

\acknowledgments
Acknowledgments.
The SEST is operated jointly by ESO and the Swedish National Facility
for Radio Astronomy, Onsala Space Observatory at Chalmers University
of Technology. MT acknowledges the support of the Academy of Finland
for our SEST observing projects.  
The work of AM is supported by the Australian Research Council.  
Part of this work was undertaken at the Jet Propulsion Laboratory,
California Institute of Technology, under contract with the National
Aeronautic and Space Administration, while SJT held an NRC-NASA/JPL
Research Associateship.  
AC is funded through a National Research Council Research Associateship.
This research has made use of the NASA/IPAC extragalactic database
(NED), which is operated by the Jet Propulsion Laboratory, California
Institute of Technology, under contract with the National Aeronautic
and Space Administration. We thank Werner Collmar for information
about the COMPTEL observations.


%-------------- REFERENCES:

\begin{thebibliography}{}
\bibitem[1998]{bertsch} Bertsch D., IAUC 6807
\bibitem[1978]{bohlin} Bohlin, R.C., Savage, B.D., Drake J.F. 1978, 
  \apj, 224, 132 
\bibitem[1975]{bolton} Bolton, J.G., Butler, P.W. 1975, AuJPA, 34, 1
\bibitem[1994]{bowen} Bowen, D.V., Osmer, S.J., Blades, J.C. et al. 
  1994, \aj, 107, 461
\bibitem[1994]{brinkmann} Brinkmann, W., Siebert, J., Boller, Th. 1994, 
  \aap, 281, 355
\bibitem[1999]{chen} Chen A. et al. 1999, \apj, in preparation
\bibitem[1998]{collmar} Collmar, W. 1998, private communication
\bibitem[1978]{condon} Condon, J.J., Jauncey, D.L., Wright, A.E. 
  1978, \aj, 83, 1036
\bibitem[1993]{dermer} Dermer, C.D. \& Schlickeiser, R. 1993, \aap, 416, 458
\bibitem[1996]{fey} Fey, A.L., Clegg, A.W., Fomalont, E.B. 1996, 
  \apjs, 105, 299 (FCF)
\bibitem[1998]{fossati} Fossati, G., Maraschi, L., Celotti A.,
  Comastri, A., Ghisellini G. 1998, \mnras, 299, 433
\bibitem[1996]{ghi:mad} Ghisellini, G. \& Madau, P. 1996, \mnras, 280, 67
\bibitem[1993]{ghisellini} Ghisellini, G., Padovani, P., Celotti, A., 
  Maraschi, L. 1993, \apj, 407, 65
\bibitem[1994]{griffith} Griffith, M.R., Wright, A.E., Burke, B.F., Ekers, 
  R.D. 1994, \apjs, 90, 179
\bibitem[1990]{imp:tap} Impey, C.D. \& Tapia, S. 1990, \apj, 354, 124
\bibitem[1974]{jones} Jones, T.W., O'Dell, S.L., Stein, W.A. 1974, \apj, 188, 
 353
\bibitem[1981]{kuehr} K\"uhr, H., Witzel, A., Pauliny-Toth, I.I.K., Nauber, U.
 1981, A\&AS, 45, 367
\bibitem[1997]{kurfess} Kurfess, J.D. 1997, in Proceedings of the
 4th Compton Symposium, ed. Dermer, C.D., Strickman, M.S. \& Kurfess, J.D.,
 AIP Conf. Proc. 410 (New York: AIP),  509
\bibitem[1973]{landolt} Landolt, A.U., 1973, \aj, 78, 959
\bibitem[1981]{large} Large, M.I., Mills, B.Y., Little, A.G. et al. 1981, 
 \mnras, 194, 693
\bibitem[1999]{macomb} Macomb D.J., Gehrels, N., \& Shrader R. 1999,
 \apj, 513, 652
\bibitem[1987]{marscher} Marscher, A.P. 1987, in Superluminal Radio Sources, 
  ed. Zensus \&  Pearson, (Cambridge University press), 280
\bibitem[1993]{mattox} Mattox, J., Bertsch, D.L., Chiang, J. et al. 1993, 
  \apj, 410, 609
\bibitem[1996]{moellenbrock} Moellenbrock, G.A., Fujisawa, K., Preston, R.A.
  et al. 1996, \aj, 111, 2174
\bibitem[1995]{montigny} von Montigny, C., Bertsch, D.L., Chiang, J. et al.
  1995, \apj, 440, 525
\bibitem[1986]{morabito} Morabito,  D.D., Niell, A.E., Preston, R.A. et al. 
  1986, \aj, 91, 1038
\bibitem[1996]{muecke} M\"ucke, A., Pohl, M., Reich, W. et al. 1996, 
  A\&AS, 120, 541
\bibitem[1990]{pks} Parkes Catalogue 1990, Australia Telescope National 
  Facility
\bibitem[1976]{peterson} Peterson et al. 1976, \apj, 207, L5
\bibitem[1993]{qui:cer} Quiniento, Z.M. \& Cersosimo, J.C. 1993, 
 A\&AS, 97, 435
\bibitem[1997]{reuter} Reuter, H.-P., Kramer, C., Sievers, A. et al. 1997,
 A\&AS, 122, 271
\bibitem[1993]{robertson} Robertson D.S., Carter, W.E., Ray, J.R. et al. 
  1993, \aj, 105, 353
\bibitem[1996]{sau:kil} Sault, B. \& Killeen, N. 1996, MIRIAD Users Guide.
\bibitem[1994]{shepherd} Shepherd, M.C., Pearson, T.J., Taylor, G.B. 1994, 
 BAAS, 26, 987
\bibitem[1994]{sikora} Sikora, M., Begelmann M.C., Rees M.J. 1994, \apj, 
 421, 153.
\bibitem[1988]{steppe88} Steppe, H., Salter, C. J., Chini, R. et al. 1988, 
 A\&AS, 75, 317
\bibitem[1992]{steppe92} Steppe, H., Lichti, S., Mauersberger, R. et al. 
 1992, A\&AS, 96, 441
\bibitem[1993]{steppe93} Steppe, H., Paubert, G., Sievers, A. et al., 1993,
 A\&AS, 102, 611
\bibitem[1994]{ter:val} Ter\"asranta H. \& Valtaoja E. 1994, \aap, 283, 51
\bibitem[1993]{tornikoski93} Tornikoski, M., Valtaoja, E., Ter\"asranta, H. 
   et al. 1993, \aj, 105, 168
\bibitem[1996]{tornikoski96} Tornikoski, M., Valtaoja, E., Ter\"asranta, H. 
   et al. 1996, A\&AS, 116, 157
\bibitem[1996]{val:ter} Valtaoja, E. \& Ter\"asranta, H. 1996, A\&AS, 120, 491
\bibitem[1998]{val:lah} Valtaoja, E., L\"ahteenm\"aki, A., Ter\"asranta, H.,
 Lainela, M. 1999, \apjs, 120. 95
\bibitem[1986]{wilkes} Wilkes, B.J. 1986, \mnras, 218, 331
\bibitem[1983]{wright} Wright, A., Ables, J.G., Allen, D.A. 1983, \mnras, 
205, 793
\bibitem[1982]{zombeck} Zombeck, M.V. 1982, Handbook of Space Astronomy and 
Astrophysics (Cambridge University Press)
\end{thebibliography}

\newpage
%
\figcaption[f1.eps] % mm_EGRET_990428.eps
{The 90 GHz and 230 GHz flux density curves for PKS~2255$-$282.
The data are from IRAM (Steppe et al. \cite{steppe88}, \cite{steppe92},
\cite{steppe93}; Reuter et al. \cite{reuter}) and SEST (this paper).
The epochs when the source was in the EGRET field of view are indicated
by dotted lines. The epoch of the EGRET-detected $\gamma$--ray outburst 
is indicated by a solid line. During the pointings from February 1993
to August 1994 PKS~2255$-$282 was $30\arcdeg$ off-axis where the sensitivity
of EGRET is very low.
\label{90_230_EGRET}}
%
\figcaption[f2.eps] % flareplot_radio_990428.eps
{ATCA and SEST observations during the millimeter-flare.
\label{flareplot_radio}}
%
\figcaption[f3.ps] % VLBI.ps
{The 5 GHz VLBA image of PKS~2255$-$282 on 1997 August 13.
\label{VLBI_image}}
%
\figcaption[f4.ps] % field.ps
{$10' \times 10'$ field showing stars measured at CTIO 
(PKS~2255$-$282 indicated by tick marks). (PS image courtesy of STScI.)
\label{field}}
% This figure is to extend over two columns, if possible:
\figcaption[f5.ps] % SED3.ps
{Spectral energy distribution in different activity states: 
quiescent state (dashed line), pre-flare state (dotted line) and flare 
state (solid line). The symbols for data points are as defined in the 
legend, and the downward pointing arrows denote upper limit values 
during the flare state.
\label{SED}}

% The table captions:
\clearpage

{\sc TABLE} 1. Average radio flux densities before the $\gamma$--ray
flare (pre-flare state).
{\sc TABLE} 2. Radio flux densities during the time of the $\gamma$--ray 
flare.
{\sc TABLE} 3. Radio flux densities after the $\gamma$--ray flare.
{\sc TABLE} 4. VLBI data.
{\sc TABLE} 5. Reference stars.
{\sc TABLE} 6. Post-flare optical observations.
{\sc TABLE} 7. EGRET and XTE-observations during the $\gamma$--ray
flare.

\end{document}


