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%
   \title{Evolution and magnetic topology \\ of the M 1.0 flare of October 22, 2002}


   \author{A.\ Berlicki
          \inst{1,2}
          \and
          B.\ Schmieder \inst{1,3}
          \and
          N.\ Vilmer \inst{1}
          \and
          G.\ Aulanier \inst{1}
          \and
          G.\ Del Zanna \inst{4}
          } 

   \offprints{A. Berlicki,\\
   \email{arkadiusz.berlicki@obspm.fr}}


   \institute {Observatoire de Paris, Section de Meudon, LESIA,
              F-92195 Meudon Principal Cedex, France
             \and
              Astronomical Institute of the Wroc{\l}aw University,
              ul. Kopernika 11, 51-622 Wroc{\l}aw, Poland
             \and
              ITA, Oslo, Norway
             \and
              DAMTP, University of Cambridge, Cambridge, UK
              }


   \titlerunning{Evolution \& magnetic topology}
   \authorrunning{A. Berlicki \etal}
   \date{Received date / Accepted date }



   \abstract{
During a coordinated observational campaign between ground based instruments
(THEMIS and VTT) and space observatories (SOHO/CDS and MDI, TRACE and RHESSI),
an M1.0 confined flare has been observed mainly during its gradual phase. The objective
of this study is to use these multiwavelength observations to study the morphology
and evolution of the flare, to analyze the dynamics
and energetics of the gradual phase and to understand the
role of various heating mechanisms. The RHESSI data allow us to
discriminate between thermal and non--thermal processes and the
RHESSI spectra indicate that the emission of the flare was mainly of thermal
origin with small non--thermal component. During the flare the RHESSI
observed emission only within 3 -- 25 keV spectral range. The temperature of plasma
obtained from the fitting of RHESSI X--ray spectra was between 8.5 and 14 MK.
The lower limit of temperature is typical for the plasma contained in
''post flare'' loops observed in X--rays. The higher temperature was
observed at about 16:25 UT during the local maximum of small impulsive
feature. The SOHO/CDS observations performed in EUV Fe XIX line also
confirm the thermal nature of the radiation. The Fe XIX emission arise at
temperatures similar to these obtained from RHESSI spectra (8 -- 9 MK).
The magnetic topology analysis of the AR from a linear force-free field
extrapolation explains the occurence of the flare due to the
reconnection of an emerging flux in a pre-existing magnetic configuration.

   \keywords{Sun: flares -- Sun: X-rays -- Sun: magnetic fields}

   }

   \maketitle

%
%________________________________________________________________


\section{Introduction}

\begin{figure*}
  \centering
  \mbox{
  \includegraphics[width=9cm]{fig_them_cont_1647.ps}
  \includegraphics[width=9cm]{themis1541_vtt1530.ps}
       }
  \caption{Image of AR 0162 obtained from THEMIS/MSDP observations
           taken between 15:41 and 16:29 UT
           in the Na I ${\rm{D_{1}}}$ spectral line wing at 200 m{\AA}
           from the line center (left panel). In the right panel we
           present the line-of-sight magnetic field deduced from
           $I$ and $V$ Stokes profiles at 100 m{\AA} from the
           Na I ${\rm{D_{1}}}$ line centre overlaid with the contours
           of the H$\alpha$ flare observed by VTT/MSDP at 15:30:01 UT.
           The field-of-view (FOV) is 400'' $\times$ 250''.}
  \label{themis_gen}
\end{figure*}


The gradual phase of flares can last half an hour to few hours.
Detailed observations of flares during their gradual
phase have been recorded by SMM and later by Yohkoh
(\v{S}vestka et al. 1982, Schmieder et al. 1996 and 1997, van Driel et al. 1997). Soft
X-ray instruments in these missions observed the formation of hot loops (10-20 MK).
These loops subsequently cool and become visible as post--flare loops in lower
temperatures (EUV and H$\alpha$ lines). The hot loops are formed at higher and higher
altitudes as the time is increasing. Such events were interpreted with models
based on the Kopp and Pneuman (1976) concept. These models are also in two dimensions
(Forbes and Malherbe 1986, Forbes and Acton 1996). After the flares, magnetic field
lines reconnect at an X point or neutral sheet in the corona.
Particles are accelerated as a consequence of the reconnection.
Plasma in the loop is heated and chromospheric evaporation occurs at the
footpoints of the loops in high density plasma region.
These models are well adapted for the interpretation of flare events in 
general, not only for eruptive two ribbon flares, but also for confined 
flares (\v{S}vestka 1986) that are due to magnetic reconnection 
at (quasi-) separatrices (D\'emoulin et al. 1997) and which do not result in 
an eruption.

During the reconnection process and the gradual phase of flares, 
SMM with its HXIS telescope could image the high energy particle plasma in four
band passes of spectra with a low spatial resolution. The Hard X-ray telescope
on--board Yohkoh had a better spatial resolution but no low energy band.  
A long debate about the long duration emission of particles started in order to
understand if the particles accelerated during the impulsive phase were
trapped at the reconnection point or if they were continuously accelerated.
To distinguish between both processes, an X-ray imager and with possible
spectroscopy analysis was proposed to built.
The Reuven Ramaty High Energy Solar
Spectroscopic Imager (RHESSI: Lin \etal 2002) was launched in 2002
and continuing to operate at the present time, are observing the solar
spectrum in the energy range 3 keV -- 17 MeV.

We had the opportunity to observe in multiwavelengths a long duration M1.0
flare during a coordinated
observational campaign between ground based instruments (THEMIS and VTT) and
space observatories (SOHO/CDS and MDI, TRACE and RHESSI).
The flare followed by several local increases of X--ray emission in its decay phase.
This flare occurs in a mainly bipolar active
region (as observed by THEMIS). Looking at SoHO/EIT and LASCO data, we noticed 
that neither coronal mass ejection nor any related coronal signature of eruptive 
events (such as transient coronal dimming and fast propagating wave-like disturbance) 
were associated to he flare. Therefore, this event was certainly a confined flare. 
We observe mainly one long bright ribbon in H$\alpha$ and with
Transition Region and Corona Explorer (TRACE) at 195 {\AA}.
To understand the magnetic configuration of the flare we proceed 
magnetic extrapolation using the {\it lfff} (linear force--free field)
code developed by D\'emoulin \etal 1997.
Long low sheared horizontal field lines were drawn, they could correspond to the
TRACE loops. RHESSI and CDS Fe XIX images show elongated emission over these elongated
field lines. This agrees with the models of gradual phase already explained above.
The three--dimensional analysis of the extrapolation allows us to have a good
idea of the topology of the flare. The analysis of the RHESSI spectra allowed us
to find if non-thermal particles can still exist during the gradual
phase of the flares.

The paper is organized as follows.
In Sect. 2 we describe the observations and their coalignment.
In the third section we discuss on the topology of the flare using {\it lfff}
extrapolation. Section 4 provides description and analysis of the RHESSI X-ray spectral
observations of the flare. Section 5 summarizes our analysis and
provides a description of our results.

\begin{figure}
  \resizebox{\hsize}{!}{\includegraphics{goes_22oct02_B.ps}}
  \caption{Time evolution of GOES X-ray flux observed during the flare of
           October 22, 2002. On the plot we marked the times of observation
           of RHESSI. THEMIS observed continuously between 15:41--17:26 UT.
           TRACE and VTT/MSDP observations were performed during whole flare.}
  \label{goes_plot}
\end{figure}

\section{Observations}

The M1.0 flare was observed in NOAA 0162 active region (AR) located about N26 E21 on
October 22, 2002. This active region was a target during the coordinated
{\it SOHO} and ground--based observational campaign in October 2002.
The active region campaign included EUV spectral observations with 
the SOHO/CDS Normal Incidence Spectrometer (NIS) (Harrison, 1995).
During the passage through solar disk between 17 and 31 October, 2002
the active region consisted of a big leading spot located
in positive magnetic field polarity and a cluster of following spots of
negative magnetic field polarity (\fig{themis_gen}).

According to the GOES 9 flux curve the flare onset was at 15:29 UT and the
1 -- 8 {\AA} X-ray flux reached its maximum at 15:36 UT (\fig{goes_plot}).
During the decay phase of this flare there were several local increases
of the X-ray flux. The most important one started at about 16:20 UT and
reached its maximum at about 16:28 UT.


\begin{figure}
  \resizebox{\hsize}{!}{\includegraphics{fov1.ps}}
  \caption{Approximate orientation of FOVs of different observations
           used in the analysis of the flare of October 22, 2002. The north is up.}
  \label{fields_of_view}
\end{figure}

\subsection{Summary of the data}

Within X-ray spectral range this flare was observed by GOES 9 and in its decay phase
by RHESSI (Lin et al., 2002). The observations of RHESSI were obtained during
the orbit between 16:03 -- 16:28 UT. During this time the X-ray emission
was detected from 3 keV to about 25 keV.

The spectro-polarimetric observations of the AR 0162 were performed with
the THEMIS telescope, operating in the MSDP (Multichannel Subtractive Double Pass
spectrograph) observing mode (Mein 2002).
On October 22, 2002 during the flare we have obtained five scans of the AR 0162.
Each scan lasted about 25 minutes. From these scans five set of data were reconstructed.
Each set of data consists of the images (600''$\times$130'') in the
Na ${\rm{D_{1}}}$ sodium line (intensity, velocity and line-of-sight magnetic field maps
at different wavelengths). Field-of-view (FOV) of the THEMIS scans was too small
to cover the whole active region so to make an image of whole AR we need to make
scans with two different FOV. Finally, from five scans we obtained two full
images (400''$\times$250'') of the AR. THEMIS observations of the AR 0162
were performed also before the flare, from about 10:18 UT.

In our analysis we also used the VTT/MSDP observations in H$\alpha$ line (Mein 1991). During
the flare between 15:30 and 16:17 UT we completed 9 scans of the MSDP entrance
window covering the whole active region with a field-of-view
of about 380''$\times$170''.

During the flare of October 22, 2002, between 15:29 and 15:41 UT, TRACE
(Handy \etal 1999) obtained high cadence (40 s) filtergrams
in the Fe XII 195 {\AA} line (with a FOV
384''$\times$384'' and a 1'' resolution), providing information on plasma at 1.5 MK.
Before and after this time the TRACE observations were
obtained with slightly lower time resolution (50 s) and larger FOV
(512''$\times$512''). The solar coordinates of the TRACE images were corrected
by using the SOHO/EIT and MDI maps.
In the EUV spectral range this flare was also observed by Coronal Diagnostic
Spectrometer (CDS) onboard the SOHO satellite.

In \fig{fields_of_view} we present the relative pointings and
fields of view of the various instruments used in our analysis.

\begin{figure}
  \resizebox{\hsize}{!}{\includegraphics{counts1.ps}}
  \caption{Time evolution of RHESSI X-ray counts rates
           observed during the flare of October 22, 2002
           in 3--6, 6--12 and 12--25 keV energy bands. The counts rates
           observed in 12--25 keV energy band were multiplied by a
           factor of 10 to make the plots more clear. Letters A, B and C
           denote the phases in the X--ray flux evolution (see Sect. 2.2).}
  \label{rhessi_plot}
\end{figure}


\subsection{X-ray observations}

Within X--ray spectral range this flare was observed by RHESSI during the
orbit between 16:03 -- 16:28 UT (\fig{rhessi_plot}). These data allowed
us to analyze the images of the flare
and the spectra. During this time the flux detected from 3 keV to 25 keV was
sufficiently high to perform the reconstruction of the images with spatial resolution of
4'' but we needed to use few minutes periods to reconstruct the images
with valuable quality. We obtained the images in standard energy bands 3--6 keV,
6--12 keV and 12--25 keV using the CLEAN algorithm with front collimators 3F--9F.
During this flare there was no attenuator in the front of the grids.

Analysis of the X--ray flux from RHESSI allowed us to concentrate
our analysis on three periods of time (\fig{rhessi_plot}). Between 16:03 and
16:13 UT the X--ray fluxes exhibited a continuous decrease (range A in \fig{rhessi_plot}).
At 16:16 UT a weak peak of the fluxes in 6--12 and 12--25 keV bands, lasting
about 2 minutes, was observed (range B). After it, from about 16:18 UT we can see
continuous increasing of X-ray flux -- range C (it reached its maximum at about 16:26 UT,
just before the passage of RHESSI through the radiative belts).

\begin{figure*}
  \centering
  \mbox{
  \includegraphics[width=5cm]{rhessi_160300_smo.ps}
  \includegraphics[width=5cm]{rhessi_161500_smo.ps}
  \includegraphics[width=5cm]{rhessi_162500_smo.ps}
       }
  \caption{Sequence of RHESSI images (224'' $\times$ 224'') reconstructed
          in 6--12 keV energy band during the gradual phase of the
          flare on October 22, 2002.}
  \label{rhessi_im}
\end{figure*}

\begin{figure*}
  \centering
  \mbox{
  \includegraphics[width=7cm]{rhessi_1603_contours_time.ps}
  \includegraphics[width=7cm]{rhessi_1625_contours_kev.ps}
       }
  \caption{Left panel: The time evolution of RHESSI X-ray structures
           observed in 6--12 keV energy band. RHESSI image
           (224'' $\times$ 224'') taken at 16:03-16:09 UT
           is overlaid with the contours of X-ray structures
           observed at 16:15--16:19 and 16:25--16:28 UT.
           Right panel: The position of RHESSI structures
           at 16:25--16:28 UT observed in three different (3--6, 6--12
           and 12--25 keV) energy band. The image shows the X--ray
           structures observed in 3--6 keV energy bands overlaid with the
           contours of the structures observed at this time in
           6--12 and 12--25 keV energy bands.}
  \label{rhessi_contours}
\end{figure*}


Between 16:03 and 16:09 UT during the decaying X-ray flux
we could reconstruct the images in 3-6 and 6-12 keV
energy bands (\fig{rhessi_im}-left panel). There was no sufficient emission at higher energies.
On these images an elongated loop--like structures are well seen. The positions of these
structures were the same on both images made in the two energy bands.
On the RHESSI X-ray flux curves (\fig{rhessi_plot}) from about 16:16 UT we can see an
increasing flux in all three energy bands. This small event lasted
2 minutes and it was not visible on the RHESSI flux curves
for higher energies. During this peak at 16:15--16:19 UT we have reconstructed the images
in 3-6 and 6-12 keV energy bands (\fig{rhessi_im}-middle panel). The X-ray emission was concentrated
in two compact sources located slightly to the south of the emission
observed at 16:03 UT (\fig{rhessi_contours} - left panel). In 12-25 keV energy band the flux was too low to
make images. After these peaks, from 16:18 UT the flux began to increase.
This increase was well visible in 3-6, 6-12 and 12-25 keV energy bands
and it reached its maximum at about 16:26 UT, just before the passage
of RHESSI through the radiative belts. In the \fig{rhessi_im} at this time we can see the
elongated loop-like structures which were located to the south of
the structures observed at 16:03 and 16:15 UT (\fig{rhessi_contours} - left panel).
The images reconstructed between 16:25 and 16:28 UT in different energy ranges
show that there were no significant shifts between the positions of X--ray
structures observed at different energies (\fig{rhessi_contours} - right panel).
The increasing
emission observed by RHESSI from about 16:18 UT was also detected
by GOES 9.


\subsection{Magnetic field observation}

On October 22, 2002 the spectro-polarimetric observations of the AR 0162 were
performed with the THEMIS telescope, operating in the MSDP spectrograph observing mode.
THEMIS is a polarization free telescope, designed especially for polarimetric
observations. The polarization analyzer provides successively the Stokes
parameters $I \pm S$ ($S=V$, $Q$ or $U$) without differential blurring due to
seeing effects. Detailed description of the specific instrumental
set up of THEMIS in the MSDP observing mode can be found in Mein (2002).

\begin{figure*}
  \centering
  \mbox{
  \includegraphics[width=8.5cm]{vtt_hcen_153001.ps}
  \includegraphics[width=8.7cm]{vtt_hcen_155952.ps}
       }
  \caption{The flare of October 22, 2002 observed
           in H$\alpha$ line by MSDP/VTT spectrograph.
           The FOV is 340'' $\times$ 200''.
           R1 and R2 denote the two main ribbons of the flare.
           LS marks the leading spot of the AR 0162.}
  \label{vtt_evol}
\end{figure*}

During the flare the observations in Na I ${\rm{D_{1}}}$
sodium line were performed only for $S=V$ and we obtained the Stokes $I$ and $V$
profiles simultaneously in the 16 MSDP channels.
By using these observations the line--of--sight (LOS) magnetic field can be
estimated from the observed Stokes profiles. The center-of-gravity
method (Semel 1967) was used to determine the magnetic field. In this method
the longitudinal component $B_{\l}$ of magnetic field is proportional to the
relative shift between the centers of $I+V$ and $I-V$ profiles (for details
see the paper of Eibe et al., 2002).

\fig{themis_gen} presents the THEMIS image of AR 0162
obtained in the Na I ${\rm{D_{1}}}$ spectral line wing at 200 m{\AA}
from the line center (left panel). In the right panel we
present the line-of-sight magnetic field deduced from
$I$ and $V$ Stokes profiles of the Na I ${\rm{D_{1}}}$ line.
The pixel size of these images is 0.25'' but the spatial resolution
of the images was limited by seeing (about 1'').
The times of the THEMIS images correspond to the beginning of the
scans and to obtain one scan it took around 25 minutes.
\fig{themis_gen} presents the image composed of two scans obtained
with different FOV.

Analysis of all THEMIS images made on October 22, 2002 showed
that there were no significant changes of LOS magnetic field
configuration during all the day.


\subsection{Coronal Diagnostic Spectrometer (CDS)}

By rastering the CDS slit (from west to east), 
monochromatic images 
and velocity maps of the field of
view \mbox{(4$^\prime \times$ 4$^\prime$)} can be constructed. 
The NIS observes the 305--635 \AA\ spectral range 
with two gratings, and a spatial resolution of 
approximately 5\arcsec. The NIS spectral range includes a large 
number of lines, but due to telemetry 
limitations, only a limited number of NIS lines
are normally  extracted and telemetred to the ground.
For more details of the CDS see e.g. 
Harrison et al. (1995) and Del Zanna et al. (2001).

\begin{figure*}
  \centering
  \mbox{
  \includegraphics[width=9cm]{trace_195_153456.ps}
  \includegraphics[width=9cm]{trace1534_themis1541.ps}
       }
  \caption{TRACE 195 {\AA} image of the AR 0162 at 15:34:56 UT. R1 and R2 denote
           the two main ribbons of the flare and LS marks the leading
           spot of the AR 0162 (left panel). In the right panel we present
           the contours of longitudinal magnetic field of the AR 0162 observed
           with THEMIS in the Na I ${\rm{D_{1}}}$
           spectral line at 15:41--16:39 UT.
           Thick and thin contours denote respectively the positive
           and negative polarity magnetic field.}
  \label{trace_them_vtt}
\end{figure*}

Two types of NIS `rasters' were used:
1) a slow (1hour 20 mins), diagnostic raster
that extracts many diagnostic
lines, 2) a fast (20 mins) raster, that 
only extracts a few lines, selected to cover a wide  range of  temperatures 
(He I 522.2 \AA, log T = 4.5; O IV 554.5 \AA, log T = 5.3; 
O V 630 \AA, log T = 5.4; Mg X 625 \AA, log T = 6.0; 
Si XII 520.8 \AA, log T = 6.3; Fe XIX 592.2 \AA,  log T = 6.9).

The two rasters used the 
4\arcsec $\times$ 240\arcsec slit and  were repeated a few times over the active region
during the day.

CDS observed NOAA 0162  during the periods
08:02--13:07 UT and  15:18--18:33 UT 
with the two types of rasters.
A relatively strong flare was observed during the 
fast raster taken during 12:43--13:07 UT, 
approximately in the same location where the 
M1 flare took place.

The M1 flare that is the subject of this paper 
was recorded by CDS with a fast raster 
which scanned the active region between 15:18 and 15:43 UT.
CDS also observed the same region with a diagnostic raster
(15:43--17:07 UT).
%
The CDS slit was located in the flare area during the peak 
X-ray emission.
It should be noted that spatially-resolved EUV 
spectroscopic observations of flares are quite rare, 
although CDS has already observed a number of them, 
as described in 
Del Zanna et al. (2002a) and  Del Zanna et al. (2002b).

New features which might help in understanding the 
complex evolution of flares are emerging from the 
CDS observations.
Del Zanna et al. (2002a) presented for the 
first time spatially-resolved maps of blue-shifts in coronal lines 
during the decaying phase of a small flare.
These blue-shifts were located  at the two ends of a small loop system
and could be interpreted as related to a slow chromospheric evaporation. 
This phenomenon is new and should not be confused with the 
well-known short-duration blue-shifts observed with SMM.

In this paper, only some results concerning these two CDS rasters
are presented. A more in-depth discussion of the CDS 
observations is deferred to a future paper.

A series of standard  corrections was applied to the raw NIS data.
These include de-biasing, flat-fielding, corrections for 
 the burn-in of the lines, cosmic ray removal.
Line intensities  were obtained 
by fitting instrument profiles, following the procedures
described in Del Zanna et al. (2001), with a 
variation that takes into account the fact that after the 
SOHO loss in 1998 the NIS lines have
asymmetric broadened profiles.
The CDS pointing has an absolute accuracy of 
about 5--10\arcsec. 
Fortunately, CDS simultaneously observes lines emitted over a 
large range of temperatures, hence CDS data 
can easily be co-aligned with other data.

\begin{figure*}
  \centering
  \mbox{
  \includegraphics[width=9cm]{vtt1559_rhessi1603_1625.ps}
  \includegraphics[width=9cm]{themis1541_rhessi1603_1615_1625.ps}
       }
  \mbox{
  \includegraphics[width=9cm,bb=58 377 535 793,clip=]{trace1604_rhessi1603.ps}
  \includegraphics[width=9cm,bb=58 377 535 793,clip=]{trace1627_rhessi1625.ps}
       }
  \caption{VTT/MSDP image of the H$\alpha$ flare in the AR 0162 at
           15:59:52 UT (left upper panel) and the longitudinal magnetic
           field deduced from THEMIS/MSDP observations (15:41--16:29 UT)
           in the Na I ${\rm{D_{1}}}$ spectral
           line at 100 m{\AA} from the line centre
           (right panel) overlaid with the contours of
           the X--ray structures observed by RHESSI at 16:03-16:09,
           16:15--16:19 and 16:25--16:28 UT (right upper panel).
           In the lower panels we present TRACE 195 {\AA} images of
           the AR 0162 (16:04:38 and 16:27:04 UT) overlaid with
           the contours of the X--ray structures observed
           by RHESSI at 16:03--16:09 and 16:25--16:28 UT.}
  \label{vtt_themis_trace_rhessi1}
\end{figure*}


\subsection{Flare ribbons and coalignments}

On the H$\alpha$ VTT/MSDP images well developed flare ribbons were
observed from 15:30 UT. The orientation of two main ribbons R1 and R2
indicates the presence of a large shear in this AR
(\fig{vtt_evol}).
The brightest parts of H$\alpha$ flare ribbons, well visible at
15:30 UT, were much weaker after 7 minutes and almost dissappeared
at about 15:59 UT. We point out that the two-ribbon shape of this 
flare is not contradictory with its confined nature, since past 
flare models (D\'emoulin et al. 1997) have shown that quasi-bipolar 
photospheric flux distributions can be associated with quasi-separatrix 
layers whose intersection with the photosphere can form only two (and not 
four) thin ribbons. 

On the TRACE sequence taken from 15:30 the first increase of emission
in the Fe XIX 195 {\AA} band was observed on the image taken at
15:32:18 UT. The first increase of GOES X-ray flux of M1.0 flare was observed at 15:29 UT.
Initially, the emission was concentrated in the compact
source well visible at 15:34:56 and 15:40:39 UT (\fig{trace_them_vtt}). Later on, from
about 15:58:55 UT the TRACE
UV emission was concentrated in small, rather low, thin loops spreaded along the axis
of the AR and in few brightenings located close to the H$\alpha$ flare emission
(\fig{vtt_themis_trace_rhessi1} - lower panels).

At the beginning of the flare the two brightest parts of H$\alpha$ ribbons were located in
the following spots group, each of them in an area of opposite polarity magnetic field
(\fig{themis_gen} -- right panel).

In this part of the AR the magnetic field gradient was very high and the initial release
of the flare energy probably took place here. This idea may be supported by the TRACE 195 {\AA}
observations. In the image taken at 15:34:56 UT the brightest UV emission was
concentrated just between two H$\alpha$ flare kernels.
In the other VTT/MSDP images taken during the gradual phase between 15:59 and
16:20 UT the H$\alpha$ flare
emission was almost no longer visible at these places but the other parts of
the flare ribbons were still bright (\fig{vtt_evol}).

A similar evolution was seen in the TRACE images. The bright UV emission seen in
15:34:56 UT almost disappeared at 15:58 UT but other UV structures were
still in emission (\fig{vtt_themis_trace_rhessi1}). \fig{trace_them_vtt} (right panel)
presents the relative position of the TRACE 195 {\AA} structures observed at 15:35:56 UT
overlaid with the contours of LOS magnetic field observed by THEMIS
between 15:41 and 16:29 UT.

The coalignment of RHESSI images with the VTT/MSDP filtergrams shows that
X-ray emission observed between 16:03 and 16:09 UT is located not on
the H$\alpha$ flare ribbons but, in projection, it
is shifted above (\fig{vtt_themis_trace_rhessi1} - upper left panel). There is no shift
between the emission observed in 3-6 and 6-12 keV energy bands. Later, the X-ray
structures observed by RHESSI at 16:25 - 16:28 UT were located (in projection)
in the same place as H$\alpha$ flare emission. In the upper right panel of
\fig{vtt_themis_trace_rhessi1} we present also the contours of RHESSI X-ray
structures observed at three times between 16:03 and 16:25 UT overlaid
on THEMIS/MSDP image of LOS magnetic field.
In the lower left panel of \fig{vtt_themis_trace_rhessi1}, the contours of RHESSI image
at 16:03-16:09 UT are overlaid on the TRACE image taken at 16:04:38 UT.
Note that X-ray
emission of the flare is located just above the bright thin loop-like
structure seen in UV \mbox{195 {\AA}}. Later, the X-ray emission observed
by RHESSI at 16:25--16:28 UT was located (in projection) just in the same place
as TRACE UV emission observed at 16:27:04 UT (\fig{vtt_themis_trace_rhessi1} - 
lower right panel).

During the flare, the Fe~XIX line of CDS was very bright
(reaching 3000 photon-events in the NIS detector), and clearly 
shows where strong heating (to at least temperatures of 8 MK)
takes place.
The peak emission, recorded between 15:35 and 
15:43 UT, occurs just above the region of 
strong mixed polarity, and has a loop-like structure,
as more clearly showed by the TRACE 195 \AA\ image
(see \fig{cds_trace_vtt_rhessi} - upper left panel). This Fe XIX emission was located
also just above strong H$\alpha$ flare ribbons observed by VTT/MSDP
at 15:30:01 UT (\fig{cds_trace_vtt_rhessi} - lower left panel).
Note that the TRACE 195 \AA\ band is dominated at the 
flare site by Fe~XXIV emission, while in general 
is multithermal, with contributions from 
plasma at temperatures as low as  700 000 K 
(see Del Zanna \& Mason, 2003)

\begin{figure*}
  \centering
  \mbox{
  \includegraphics[width=8cm,bb=58 377 535 793,clip=]{TRACE1534_CDS1518_FE.PS}
  \includegraphics[width=8cm,bb=58 377 535 793,clip=]{TRACE1614_CDS1543_FE.PS}
       }
  \mbox{
  \includegraphics[width=8cm,bb=54 300 535 699,clip=]{VTT1530_CDS1518_FE.PS}
  \includegraphics[width=8cm]{HSI1625_CDS1543_FE.PS}
       }
  \caption{TRACE 195 {\AA} images of the AR 0162 (15:34:56 and
           16:14:45 UT) overlaid with
           the contours of the UV structures observed
           by CDS at 15:18:16--15:43:16 and 15:43:30--17:07:33 UT
           in the Fe XIX line (592.2 {\AA}) (upper panels) and
           VTT/MSDP image of the H$\alpha$ flare in the AR 0162 at
           15:30:01 UT overlaid with the contours of the UV
           structures observed by CDS (Fe XIX line, 592.2 {\AA}) at
           15:18:16--15:43:16 UT (left lower panel)
           and RHESSI image (245'' $\times$ 245'') reconstructed
           for the time 16:25--16:28 UT overlaid with CDS Fe XIX structures
           observed at 15:43:30--17:07:33 UT (right lower panel).}
  \label{cds_trace_vtt_rhessi}
  \end{figure*}

During the decay phase, the 
Fe~XIX emission becomes more diffuse, a common 
feature already noted in 
Del Zanna et al. (2002a, 2002b) and shown 
in \fig{cds_trace_vtt_rhessi} (right panels)
The Fe~XIX emission was recorded while the 
slit was covering the eastern half of the 
FOV, between 16:20 and 16:51 UT.

\begin{figure*}
\begin{center}
  \includegraphics[width=17cm,bb=20 255 586 745,clip=]{Figure.ps}
% \psfig{figure=Figure.ps,width=17.0cm,bbllx=1pt,bblly=260pt,bburx=580pt,bbury=800pt,clip=}
\end{center}
\caption{
({\em top left:})      superposition of the THEMIS/MSDP Na D1 magnetogram 
                       of 12:32:02 UT with a sample of magnetic field lines
                       resulting from the {\it lfff} extrapolation. 
({\em bottom right:})  partial view of the flare as observed in H$\alpha$ 
                       with the VTT/MSDP at 15:59:52 UT. the field lines 
                       are the ones which are rooted in the flare ribbons 
                       R1 and R2 as well as in the emreging bipole. 
({\em bottom left:})   same as {\em bottom right}, except that the underlying 
                       image is the THEMIS magnetogram.
({\em top right:}      projection view of the field lines shown in the 
                       {\em bottom} pannels. Continuous (resp. dashed) 
                       contours stand for positive (resp. negative) 
                       vertical fields. North is toward the left. 
\label{extrapol}
}
\end{figure*}


\section{Interpretation of the flare in a simple topology}

In this section we provide an interpretation for the triggering 
and the development of the confined flare by the means of 
a topological analysis of the coronal magnetic field above 
NOAA 0162 calculated from a linear force-free field model. 

\subsection{Linear force-free field extrapolation}

For the extrapolation, the line-of-sight magnetic fields measured 
with the THEMIS/MSDP in the Na I ${\rm{D_{1}}}$ line at 12:32:02 UT were assumed 
to be equal to the vertical fields at the altitude $z=0$ (i.e. 
$\vec{B}(x,y,z=0)$) divided by the projection angle, as often done 
in extrapolations using longitudinal magnetograms only (see e.g. 
D\'emoulin \etal 1996, Eibe \etal 2002). The extrapolation 
was performed in the linear force-free field approximation 
($\vec{\nabla} \times \vec{B} = \alpha \vec{B}$ with $\alpha=$ cst) 
with the FFT method, so with periodic boundary conditions in $(x,y)$ 
(Alissandrakis, 1981). In order to minimize the boundary effects, 
a large computational box was chosen, with $\Delta x = \Delta y = 
10^3$ Mm, and the magnetic field was computed up to $z=150$ Mm 
above the photosphere. The calculation was done with $nx=ny=1024$ 
mesh points so as to sufficiently resolve the emerging bipole in 
the trailing part of the active region. The force-free parameter 
$alpha$ was iteratively tuned so as the obtain a general fit 
between the orientations of the flare loops observed by TRACE 
and of the calculated field lines rooted in the vicinity of 
these loops. This constraint led to fix $\alpha = -5 \times 
10^{-3}$ Mm$^{-1}$, whose sign is compatible with the 
hemispheric helicity rules found by Pevtsov \etal (1995). 

Even though the results of this calculation were used to study 
the triggering and the development of the flare, we must point 
out that they are uncertain for two reasons. Firstly, the selected 
$\alpha$ is equal to $80 \%$ of the resonnant value well 
known for extrapolations performed with a FFT in a large 
computational box. This leads to a very probable artificial 
over-shearing of the large scale loops. Secondly, it was 
impossible to find a single value of $\alpha$ with, even 
in a $20 \%$ interval, which allowed to fit all the observed 
TRACE loops with calculated magnetic field lines. Therefore,
it appears that the $\alpha=$cst hypothesis of the {\it lfff} 
extrapolation may be strongly disputed in this region. 
Nevertheless, as the global magnetic topology is a well 
preserved property of magnetic fields since it is an intergral 
quantity, we still used the {\it lfff} model to pursue 
our analysis, although it is clear that it is disputable 
on small scales (where electric currents are concentrated) 
and on large scales (where the field is too much sheared).

\begin{figure*}
  \centering
  \mbox{
  \includegraphics[width=9cm]{spectrum_1603_1604.ps}
  \includegraphics[width=9cm]{spectrum_1613_1614.ps}
       }
  \mbox{
  \includegraphics[width=9cm]{spectrum_1616_1617.ps}
  \includegraphics[width=9cm]{spectrum_1625_1626.ps}
       }
  \caption{RHESSI spectra reconstructed during the gradual phase of the
           solar flare in four times: 16:03--16:04, 16:13--16:14, 16:16--16:17
           and 16:25--16:26 UT. Dotted lines denote the calculated thermal component,
           dashed lines represent non--thermal component, thin continuous lines
           are the Fe--line feature ($\sim 6.7$ keV) calculated in our model. In our modelling
           of the specrum at 16:25--16:26 UT we also used Fe/Ni--line feature ($\sim 7.9$ keV).
           Thick continuous lines are the sum of all components -- predicted
           photon spectum. With crosses we marked observed flare spectrum.}
  \label{rhessi_spectra}
  \end{figure*}

\begin{figure*}
  \centering
  \mbox{
  \includegraphics[width=9cm]{RHESSI_T.PS}
  \includegraphics[width=9cm]{RHESSI_EM.PS}
       }
  \caption{The time evolution of the temperature $T$ and emission
           measure $EM$ of the flare on October 22, 2002 These
           parameters were obtained from the fitting the observed
           RHESSI spectra with the F-VTHC-BPOW-NLINE model.}
  \label{temp_emis}
  \end{figure*}


\subsection{Origins of the asymmetrical ribbons}

Figure \fig{extrapol} shows the results of the {\it lfff} 
extrapolation. The large scale shear is evident on the 
scale of the full active region, as seen on the {\em top 
left} pannel. The other pannels show the connectivies 
of the magnetic field lines in the vicinity of the emerging 
bipole and around the flare ribbons as observed in H$\alpha$. 

In this region, even though a few nulls points are present 
around the negative moving magnetic features on the edges 
of the leading spot, the global topology is {\em not} complex : 
no thin quasi-separatrix layers were found to be associated 
with the flare loops, contrary to many previous studies of 
confined flares (see e.g. Mandrini \etal 1996,  Schmieder 
\etal 1997, Aulanier \etal 2000). Also, the connectivities 
look very continuous around the region of the emerging bipole. 
So at first sight it appears that the flare studied in this 
paper cannot be interpreted with the standar reconnexion model. 
However emerging bipoles are known to often emerge twisted 
(Leka \etal 1996, Lopez-Fuentes \etal 2000), which is 
supported by the Huairou vector magnetograms for this 
region (Y.H. Yan, privated communication). Therefore 
we conjecture that the flare was still probably due to 
magnetic reconnexion, but here between small scale twisted 
emerging fields (not calculated in the extrapolation) and the 
large scale sheared magnetic fields of the global bipole of 
NOAA 0162. Since the flare was confined, it may then be 
interpreted by the global disruption model calculated by 
Amari et al. (1999) in a very symetrical configuration, 
where a very locally twisted flux tube undergoes an ideal 
kink instability and reconnects with overlaying potential 
fields. 

If we assume the global disruption scenario, the {\it lfff} 
model naturally provides an explanation for the appearance 
of the compact ribbon R1 and its elongated counterpart R2, 
as well as its fast propagation away from the emerging 
bipole and toward the leading sunspot. Figure \ref{extrapol}
indeed shows that the field lines located higher and higher 
above the emerging bipole are rooted in R1 and R2, and have 
footpoint separations that are small on the R1 side and 
quite large on the R2 side. The reconnexion of the growing 
emerging fields with higher and higher overlaying fields 
sequentially provides a natural explanation for 
the sequential brightening of their footpoints along 
the flare ribbons, included along the very elongated 
R2 ribbon. In this frame, the fast propagation of the 
brightenings as observed by TRACE are not due to MHD 
waves, but are rather due to the electrons beams accelerated 
from the reconnexion regions above the emerging bipole 
in the long overlaying loops of NOAA 0162, which progressively 
reach larger and larger chromospheric areas toward the West. 

It is also noteworthy that the shape of the field lines 
anchored in the H$\alpha$ ribbons follows the elongated 
shape of the RHESSI images recorded in the $6-12$ keV 
range of energy between 16:03-16:09 UT. Thus the extrapolation 
supports the idea that RHESSI permits to see thermal emission 
of coronal loops heated by the confined flare, though it 
is not contradictory with a possible non-thermal component 
since its is known that reconnexion can accelerate particules 
along reconnected field lines. The RHESSI emission at 16:25-16:28 
UT being similar in shape to the one recorded at 16:03-16:09 UT, 
though slightly shifted southward, suggests that this secondary 
event may be associated to the resistive relaxation of the system 
which may have been forced too rapidly by the flare of 15:59 UT, 
as in the reconnexion calculations of Karpen et al. (1998). 
Unfortunately, the present data and model cannot provide more 
clue to this issue. 


\section{RHESSI X--ray spectra of the flare}

Using the RHESSI data and SPEX package we also performed the analysis
of spectra obtained from the data taken between 16:03 and 16:27 UT.
To substract the background emission we used the data taken before the
flare during the previous RHESSI orbit at 15:12--15:24 UT. Then, we
made the spectra for all 1 minute time bins between 16:03 and 16:27 UT.
For the spectra we used modified binning code No. 14. To increase the
spectral resolution of spectra, instead of 1 keV energy bins, between
3 and 10 keV we have used 0.25 keV energy bins. To make a spectra we
used the collimators 1F -- 9F, except 2F and 7F. We omitted the data from
these detectors because they have a poorer energy resolution than
intended (Smith \etal 2002). We performed the
full calculations of all diagonal and off-diagonal matrix elements.

Using the procedures of SPEX package we have reconstructed the spectra
for different periods in time. All the spectra made between 16:03 and
16:27 UT exhibited similar distribution of the energy from 3 to 10 keV.
In this range we could fit them by using thermal model of the
SPEX package. On the spectra plotted both in counts
and photon scale the spectral line feature on around 6.7 keV was well
seen (\fig{rhessi_spectra}). 
Above about 10 keV the observed spectra did not follow the thermal
model and we could fit this part of the spectra using non-thermal model.
Finally, we performed the fitting of the spectra contained within
the energy range 3.0 -- 20 keV using F-VTHC-BPOW-NLINE model. This model is a mixture
of thermal and non-thermal components together with the lines of given intensity,
wavelength and width.
In our fitting procedure we have used the
line on around 6.7 keV which describe the Fe--line feature emission (Phillips 2003).
For the spectra reconstructed during the increase of X--ray flux at about
16:26 UT we also introduce another line at the energy of about 8 keV
which is referred to as the Fe/Ni--line feature (Phillips 2003).
These two line features are emitted from high plasma temperatures and
can be used as temperature indicators.
The 6.7 keV feature corresponds mainly to the Fe XXV
lines and associated dielectronic satellites, centered on 6.6 keV
and to Fe XXVI lines and associated satelites, centered on 6.9 keV.
The 8 keV line feature is due to Fe XXV and Fe XXVI lines as well as Ni XXVII
and Ni XXVIII lines. The lines are excited at electron $T_{e} > 10$ MK
The contributions of individual lines to the total emission within both
line features is temperature dependent.
Fe XXV line on 6.696 keV give the most intense input to Fe--line
feature and the major part of Fe/Ni--line feature coming from
Fe XXV and Fe XXVI lines. The contribution of all Ni XXVII and
Ni XXVIII lines and their satellites is always below about 25\%
of the total Fe/Ni--line feature emission.
The Fe/Ni--line feature starts to be visible at higher temperatures
than Fe--line feature and for the $T_{e} < 20 MK$, the
Fe/Ni--line feature is very weak.
The intensity ratio of Fe and Fe/Ni--line features depends on
the temperature of plasma.

For all reconstructed spectra we could find the Fe--line feature
on about 6.7 keV (\fig{rhessi_spectra}). From 16:25 UT, during
small increase of X--ray flux, the Fe/Ni--line feature appeared
on about 7.9 keV. This emission was very weak and it could
suggest for higher temperature.

The continuous spectra made between 16:03 and 16:27 UT
differs mainly in the contribution of non-thermal radiation
to the spectrum above the 10 keV energy. In the period 16:03 -- 16:13 UT
the spectra were more thermal in comparison with the spectra made
during the peak of X--ray flux at 16:16 -- 16:18 UT or during
the increasing of the radiation observed from about 16:18 UT. The non-thermal
component gave the largest input during the maximum of X-ray flux
at about 16:26 UT (\fig{rhessi_spectra}). All spectra are similar and
characteristic for the gradual phase of the flare where the non--thermal
component has a very low input.

In the \fig{temp_emis} we can see the time evolution of the temperature $T$
and emission measure $EM$ obtained by the fitting of the observed
spectra with the model F-VTHC-BPOW-NLINE.
First, the observed temperature decreases from about \mbox{9.7 MK} to \mbox{8.8 MK}
and after, at \mbox{16:10 UT} starts to increase to the value of \mbox{13.6 MK}.
It means that from about \mbox{16:10 UT} some heating processes took place.
The maximum value of the temperature at 16:25 UT is consistent with
the appearance of the Fe/Ni--line feature at the same time.
At \mbox{16:09 UT} the emission measure started to decrease from about
$2.53 \times 10^{48} cm^{-3}$ to $2.35 \times 10^{48} cm^{-3}$.


\section{Conclusions}

The multiwavelength observations of an M1.0 flare on October 22, 2002
obtained during its gradual phase allowed us to perform a complex
analysis of the morphology and to understand the mechanisms
which are responsible for the observed evolution of this event. No CME 
and no coronal signature for an eruption was observed at the time 
of this flare, which implies that this event was of 
the confined type. The magnetic field (photospheric and extrapolated) 
and the aligned H$\alpha$ ribbons shows that this active region is
highly sheared, thus stores energy. The triggering of H$\alpha$ flare
occurs in the region of high gradient of B which could correspond
to the emergence of a locally very twisted magnetic flux.
The {\it lfff} extrapolation of the magnetic field provides an
explanation for the appearance two H$\alpha$ flare ribbons:
ribbon R1 was compact and the its position was almost
constant during the flare; ribon R2 was elongated and exhibited
fast propagation from the emerging bipole towards the leading
sunspot.

H$\alpha$  and TRACE observations show the development of two
bright ribbons going away from each other versus time.
This agrees with the model already accepted for gradual phase
of flare (Forbes and Malherbe, 1986, Forbes and Acton, 1996).
During the decay phase, X-ray was observed above the post-flare loops
at the reconnection site of magnetic field lines. This emission
originated from trapped particles or continuous reconnection
processes. There was no
X-ray emission at the footpoints of the loops which
indicate that there was no non--thermal heating of H$\alpha$ flare ribbons.
The thermal conduction heating can be responsible for the TRACE ribbons
but is not sufficient to explain the H$\alpha$ flare emission
(Berlicki and Heinzel 2004). We need to investigate other processes
for instance the effect of high coronal pressure.

The RHESSI observations performed during the gradual phase
of this flare suggests that the X--ray emission observed
above the active region had a thermal nature. The RHESSI spectra
showed that the contribution of non--thermal emission to the total
X--ray flux is not strong. From the fitting of RHESSI spectra
we found the temperature of flare plasma which is not
high (around 9--14 MK) and typical for the structures observed
during the gradual phase of solar flares. The CDS observations
in Fe XIX line (592.2 {\AA}) performed at 15:43:30--17:07:33 UT
confirm the thermal nature of the radiation. The CDS Fe XIX emission
arise at the temperature around 8--9 MK which is in
agreement with the temperatures obtained from the RHESSI spectra.
Besides, the EUV emission was located in the projection in the
same place as RHESSI X-ray emission.
The intensity of X--ray lines used in fitting the RHESSI spectra
support the thermal nature of X-ray emission. The Fe--line feature
($\sim$ 6.7 keV) indicate the presence of high temperature
plasma (about 9--10 MK) (Phillips, 2003). At 16:26 UT, during the local increase
of X--ray flux we also observed another line--feature (Fe/Ni, $\sim7.9$ keV).
This line emission resulting from the increase of the temperature which
rise to about 14 MK.



\begin{acknowledgements}

The magnetic field extrapolations used in this paper have been obtained
from the code base FRench Online MAGnetic Extrapolations (FROMAGE).
FROMAGE is a joint project between LESIA (Observatoire de Paris),
CPhT (Ecole Polytechnique) and the Centre National d'Etudes
Spatiales (CNES).
The work of AB was supported by the European Commission 
through the RTN programme (European Solar Magnetism Network,
contract HPRN-CT-2002-00313).
VTT/MSDP observations have been provided and reduced by P., N. Mein
and J. Staiger. These observations were performed during a MEDOC
campaign (JOP 157). GDZ acknowledges support from PPARC (UK).
The TRACE analysis at SAO is supported by a contract
from Lockheed Martin. 

\end{acknowledgements}


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\end{document}


