\documentclass{aa}
%\documentclass[referee]{aa}

%\usepackage{txfonts}

\usepackage{epsfig}
%\usepackage{longtable}

\usepackage{natbib}
\bibpunct{(}{)}{;}{a}{}{,} %for A&A style

\def\orb[#1 #2]{{$#1^{#2}$}}
\def\tm[#1 #2 #3]{{$^{#1}{#2}_{#3}$}}

\newcommand{\lo}[1]{\mbox{$\rm _{#1}$}}
\newcommand{\up}[1]{\mbox{$\rm ^{#1}$}}
\newcommand{\beq}{\begin{equation}}
\newcommand{\eeq}{\end{equation}}


%|  Ions, in format like Ca II and Ca^+:
\newcommand{\ionr}[2]{\mbox{\ion{#1}{#2}}}
%\newcommand{\ionc}[2]{\mbox{{#1}}$^{{#2}}$}
\newcommand{\ionc}[2]{\mbox{\element[#2]{#1}}}

%\newcommand\gion[2]{#1$\;${\small\rmfamily\@Roman{#2}}\relax}% 

%|  Atomic states in LS coupling
\newcommand{\LS}[2]{\mbox{$\,^{#1}${#2}}}
\newcommand{\LSJ}[3]{\mbox{$\,^{#1}${#2}$_{#3}$}}
\newcommand{\LSJp}[4]{\mbox{$\,^{#1}${#2}$_{#3}^\textrm{{#4}}$}}

\newcommand{\specline}[3]{\ensuremath{\ion{#1}{#2}\; \lambda{#3}}}


\newcommand{\sixii}{\specline{Si}{xii}{520.665}~\AA}

\newcommand{\fexix}{\specline{Fe}{xix}{592.16}~\AA}


\newcommand{\SOHO}{\mbox{SOHO}}       % SOlar Heliospheric Observatory
\newcommand{\CDS}{\mbox{CDS}}         % Coronal Diagnostic Spectrometer
\newcommand{\NIS}{\mbox{NIS}}         % Normal Incidence Spectrometer
\newcommand{\GIS}{\mbox{GIS}}         % Grazing Incidence Spectrometer
\newcommand{\SUMER}{\mbox{SUMER}}     % Solar UV Measurement of Emitted Radiation
\newcommand{\SERTS}{\mbox{SERTS}}     %  Goddard Solar Extreme
                                %  Ultraviolet Research Telescope and Spectrograph

\newcommand{\UVCS}{\mbox{UVCS}}

\newcommand{\Skylab}{\mbox{Skylab}}   % ...
\newcommand{\YOHKOH}{\mbox{YOHKOH}}   % ...
\newcommand{\SXT}{\mbox{SXT}}         % " Soft X-ray Telescope
\newcommand{\CHIANTI}{\mbox{CHIANTI}} % ...



%|  Electron density et al.
\newcommand{\elN}{\ensuremath{n_\mathrm{e}}}
\newcommand{\elT}{\ensuremath{T_\mathrm{e}}}
\newcommand{\genT}{\ensuremath{T}}
\newcommand{\elP}{\ensuremath{P_\mathrm{e}}}

\newcommand{\Ne}{$N_{\rm e}$}
\newcommand{\Te}{$T_{\rm e}$}

\newcommand{\cmt}{cm$^{-3}$}
\newcommand{\cm}{cm$^{-1}$}

\def\rsun{R$_\odot$}

\newcommand{\fig}[1]{Fig.~\ref{#1}}
\newcommand{\etal}{et~al. }
\newcommand{\tab}[1]{Table~\ref{#1}}
\newcommand{\eq}[1]{Eq.\,(\ref{#1})}


%natbib:
%\cite
%\citet me (1999)
%\citep (me 1999)
%\citealt me 1999
%\citeauthor me
%\citeyear 1999

\begin{document}

\title{Spectroscopic signatures of the M 1.0 flare on October 22, 2002}

\author{G. Del Zanna
\inst{1}\thanks{Previous address: DAMTP, University of Cambridge} 
\and
A.\ Berlicki \inst{2,3} 
\and
H. E. Mason
\inst{4}  
\and
          B.\ Schmieder  \inst{2,5}
}
\institute{
MSSL, University College London Holmbury St. Mary Dorking UK RH5 6NT
\and
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
Department of Applied Mathematics and Theoretical Physics,
University of Cambridge, Cambridge UK  CB3 0WA
   \and
   Institute of Theoretical Astrophysics, University of Oslo, 
   Blindern, N-0315 Oslo, Norway
}


\offprints{G. Del Zanna \\ (G.Del-Zanna@damtp.cam.ac.uk)}

\date{\today}

\titlerunning{The M~1 flare on October 22, 2002 }

\abstract{
In this paper we present a further study of the M1 class flare
observed on 22nd October 2002. We focus on the SOHO Coronal 
Diagnostic Spectrometer (CDS) spectral observations
and the ground-based (VTT/MSDP) observations 
performed during a multi-wavelength campaign with TRACE.
Strong blue-shifts are observed in the CDS coronal and
flare lines during the impulsive phase of this flare.
From a careful wavelength calibration  
we deduce upflows of 140 km/s for the Fe~XIX emission.
We relate these to upflows deduced from the H$\alpha$ features seen
with the VTT/MSDP. The H$\alpha$ emission is analysed using non-LTE models.
These results contrast with other publications which
show downflows in H$\alpha$ emission correlated with upflows
in high temperature emission.
We interpret our results as due to 'gentle evaporation'.
In the upflow regions we measure electron densities 
$\simeq$10$^{11}$ cm$^{-3}$ in transition region lines.
SOHO/MDI magnetograms show changes during the 
flare, indicative of flux cancellation.
We discuss the importance of spectral observations in relation to
our understanding of the development of flares. We consider the complexity
and limitations of current spectroscopic observations and look towards
future solar missions, such as Solar-B.
%
\keywords{Sun: corona -- Techniques: spectroscopic -- Line: profiles}
}

\maketitle



\section{ Introduction }

An important aspect of flare evolution is the presence of flows.  
These are characterised by asymmetrical profiles of spectral lines. 
These asymmetries or line shifts (particularly in the blue wing)
have been observed both 
in X-ray and in chromospheric lines. They are 
thought to be caused predominantly by vertical mass motions 
in flaring layers of the solar atmosphere. It is now widely accepted that 
during  flares the chromospheric plasma is evaporated 
into the corona providing material for
the so-called post-flare loops (\citealt{antiochos_sturrock:78}). 

The broadening and blue-shifting of spectral lines for high temperature
emission during the initial phase of solar flares has been
extensively studied in the X-ray wavelength range
using data from SOLFLEX (cf. \citealt{doschek_etal:79}), the Solar Maximum
Mission's (SMM's) X-ray polychromator (XRP) 
 Bent Crystal Spectrometer
(BCS) (cf \citealt{antonucci_etal:82}) and the BCS on YOHKOH
(see, e.g., \citealt{bentley_etal:94}). Plasma velocities deduced from
Doppler-shift analysis were of the order of a few hundreds 
$\mathrm {km\;s^{-1}}$. This blueshifted emission was interpreted as
evidence for chromospheric evaporation. However,
these BCSs had no spatial resolution, so it was not possible 
to locate the source 
of the blueshifted X-ray emission, or to spatially separate this 
from any stationary components. 
A further complexity in the interpretation of BCS spectra 
was due to the superposition (at different wavelengths)
 of stationary components originating from different locations.
In general asymetric line profiles
were observed, rather than totally shifted lines.
Without spatial resolution in the BCS X-ray observations, 
any weak blue-shifted X-ray emission would swiftly be swamped by 
the strong stationary emission from the flare X-ray loops.
In a  recent paper, \cite{warren_doschek:05}
interpret these blue-shifted asymetries as a succession 
of independently heated
threads within flare loop structures. To progress further, it
is essential to spatially locate the high temperature 
blue shifted emission in relation to the flare loops and structures.
This can only be achieved with a spectrometer which also has 
good spatial resolution. Observations of this kind are extremely rare.

The SMM XRP Flat Crystal Spectrometer (FCS) had a limited
spatial resolution (14$^{\prime\prime}$). \cite{zarro_etal:88} 
reported spatially confined Mg XI blue-shifted 
emission during the impulsive phase of a small flare.
The Fe~XXI line at 1354.1~\AA\, emitted at
about $10^7$~K, was studied with the SMM Ultraviolet
Spectrometer and Polarimeter (UVSP) instrument (with spatial 
resolution down to 3$^{\prime\prime}$) . 
\cite{mason_etal:86} deduced upward flows
in Fe~XXI emitting plasma in excess of 200~km/s.
Their observations supported the hypothesis of 
chromospheric evaporation in the
'footpoint region' of what subsequently became soft X-ray emitting loops.
\cite{schmieder_etal:90} studied the evaporation in the 
transition region during the gradual phase of solar flares.
From SMM/UVSP CIV emission and HXIS data, they found both upflows
of 4-12 $\mathrm{km\;s^{-1}}$ and downflows of about 10 $\mathrm{km\;s^{-1}}$. 
They interpreted the results in terms of the chromospheric ablation 
model proposed by \cite{forbes_etal:89}.

The Coronal Diagnostic Spectrometer (CDS) instrument 
on SOHO has the capability to simultaneously 
obtain spatial and spectral information for emission over a 
wide range of plasma temperatures. One of the  hottest lines recorded
with CDS is the Fe~XIX flare line at 592.2~\AA. This is formed at a
temperature around $8 \times 10^6$~K. Using CDS, it is possible
to locate the source of high temperature
blue-shifted emission seen during a solar flare. 
\cite{czaykowska_etal:99} reported SOHO/CDS 
observations of blue-shifts in Fe XVI 
($3 \times 10^6$~K) and Fe~XIX ($8 \times 10^6$~K) lines.
These observations were however performed 
more than one hour after the maximum 
of a two-ribbon M6.8 flare. The blue-shifts appeared to be 
located at the outer edges of the flare ribbons.

Spatially-resolved spectral information from CDS were 
reported by \cite{delzanna02_cospar}
covering the entire evolution of a small flare.
Long-lasting (over an hour) blue-shifts of 
30~km/s in coronal lines (Mg~X, Si~XII, emitted at 
1-2~MK) were located at the footpoints of a small loop system.
These could be interpreted as indirect evidence of chromospheric evaporation. 
The hotter Fe~XIX emission also showed some
interesting properties, including a pre-flare enhancement 
along the loop structure (\citealt{gibson_etal:02}).
A long-term campaign of SOHO/CDS observations over 
2001-2004 has provided evidence that similar 
features are common in C-class flares 
(see \citealt{del_zanna_etal_prague}).

Recently, \cite{teriaca_etal:03} published
CDS and ground-based observations of 
an eruptive C2 flare.
The main findings were quasi-simultaneous and co-spatial 
measurements of upflows in the flare line (Fe~XIX) 
and downflows in the chromosphere, as seen in the 
He~I 10830~\AA\ line. 
\cite{brosius_phillips:04} presented CDS observations
of a larger M6 flare.
These authors kept the position of the CDS slit fixed,
thus sacrifying spatial coverage in favour of 
a high (10s) time resolution. They used TRACE observations
to track the evolution of the flare.
Their main results were measurements of upflows ($\simeq$60 km/s) 
in flare lines and downflows ($\simeq$40 km/s) in cooler
transition region lines during the impulsive phase.
The X-ray emission appeared $\simeq$1m after the start of 
the impulsive intensity rise in the cooler lines.
They concluded that their observations were consistent with 
non-thermal particles 
transferring most of the energy into the chromosphere and 
creating these oppositely-directed flows. 

The ''explosive'' evaporation seen in the emission from high
temperature flare plasma is thought to be driven 
by non-thermal electrons accelerated during the primary energy release
(\citealt{antonucci_etal:84}) and may occur when the heating rate due 
to collisions of non-thermal 
electrons with the ambient plasma in upper layers of the 
chromosphere exceeds the chromospheric
radiative losses (\citealt{fisher:87}). This occurs 
when the temperature of the heated region is 
above \mbox{$2 \times 10^{5}$ K} and thus this kind of 
evaporation cannot be seen in chromospheric spectral lines which are
formed in a cooler plasma.

If a region of the chromosphere heated by non-thermal electrons 
is thick enough, then the 
rapid temperature increase produces an enhanced pressure 
in the heated region. This overpressure, 
besides the evaporation, also drives downward-moving cool and dense 
{\it chromospheric condensations} (\citealt{fisher_etal:85c}) which seem 
to be responsible for red-shifts 
of the H$\alpha$ line profiles reported by many authors 
(\citealt{svestka:76}, \citealt{ichimoto_kurokawa:84}).
\cite{fisher_etal:85a} modelled the hydrodynamic 
and radiative response of the atmosphere to short impulsive injections 
of non-thermal electron beams. They showed that only the high-energy 
flux of non-thermal electrons drives explosive evaporation, 
accompanied by the formation of 
cool chromospheric condensations in the flare chromosphere.

Another situation occurs when the flux associated with 
non-thermal electrons is low. 
Then only a weak chromospheric evaporation takes place. 
This kind of evaporation
is referred to as {\it gentle evaporation} 
(\citealt{antiochos_sturrock:78}; 
\cite{fisher_etal:85b};
\citealt{schmieder_etal:87})
and it can be observed in transition region lines and in 
chromospheric spectral lines like H$\alpha$ or 
in Ca II (8542 {\AA}). Antiochos \& Sturrock suggested that the 
gentle chromospheric evaporation may occur after the primary energy 
release when the non-thermal 
electron flux has stopped. This evaporation could be driven 
by the large conductive 
heat flux from a high-temperature flare plasma contained 
in magnetic tubes above the 
chromosphere
(\citealt{fisher_etal:85b}).
Such physical conditions may appear during the 
gradual phase of solar 
flares, when there is no significant flux of non-thermal electrons. 
In the \cite{forbes_etal:89} model for flare-loop formation 
by magnetic reconnection 
the conduction of the thermal energy generated at 
the slow-mode shocks drives a 
gentle evaporative upflow from the ribbons. 
%

\cite{schmieder_etal:87}
observed small but long-lasting blue shifts in flare ribbons 
in the H$\alpha$ line during the gradual phase 
of three solar flares and interpreted them as due to 
upflows with velocities $<\;10\; \mathrm{km\;s^{-1}}$. 
These upflows were thought to be caused by a gentle 
chromospheric evaporation driven by the heat 
conduction along the field lines connecting the chromosphere 
with a reconnection site in the corona. 
In the case of gentle evaporation, the observed flow velocities are
much lower than during the impulsive phase of flares.


\begin{figure*}[!htbp]
\centerline{\epsfig{file=goes3_bw.ps, width=10.0cm , angle=90 }}
%  \resizebox{\hsize}{!} {\includegraphics{goes2_bw.ps}}
\caption{Time evolution of the GOES 8  X-ray fluxes
on October 22, 2002 during the M1 flare and the CDS raster 
scan.
The dot-dashed curve is the time derivative of the 
0.5--3 \AA\ light curve, and can be used as a proxy for the 
hard X-ray emission. 
 The timings of a few 
key observations are indicated 
(top: CDS slit positions, see 
Fig.~\ref{fig:trace_seq}; middle: TRACE images, see Fig.~\ref{fig:trace_seq_zoom};
bottom: H$\alpha$ slit positions during two 
consecutive rasters, see  Fig.~\ref{fig:vtt}).
}
  \label{fig:goes_plot}
% Fig.~\ref{fig:goes_plot}
\end{figure*}




The analysis of 
gentle evaporation during the gradual phase using chromospheric lines
is quite complex since the line profiles can be in absorption or
enhanced emission. 
Very often we also observe
self-reversed profiles. Deducing the velocity from 
such profiles is rather difficult and cannot be done 
just by searching for Doppler shifts 
with the bisector method
 (\citealt{heinzel_etal:94}).
Chromospheric line profiles suggest that 
the flare atmosphere is highly dynamic and 
stratified with rather complicated plasma 
motions. 
In addition,  these lines are optically thick.
The most reliable way to analyse the velocity 
field is to use the non-LTE  radiative transfer codes,
which enable us to compute the chromospheric models 
with velocity fields. Resulting synthetic 
line profiles can then be compared with the observed ones. 
The analysis by  \cite{teriaca_etal:03} is severely
limited in this respect, since the velocities in the chromosphere were
obtained soley from the relative shifts between Gaussian fitted profiles.
This can lead to inaccurate results concerning the magnitude and even
the direction of the chromospheric flows.

A coordinated observational campaign
was organised to follow the evolution of the 
NOAA 0162 active region during October 22, 2002.
Ground based instruments (THEMIS and Vacuum Tower Telescope - VTT - 
on Tenerife, Canary islands, Spain) and
space observatories (SOHO, TRACE and RHESSI)
were involved. 
The active region had several flares, including an M1 flare.
The morphology and evolution of the M1 
flare in particular during the gradual phase 
(when RHESSI was observing) were presented in \cite{berlicki_etal:04}.
Here, we focus in particular on the nature of the emission from
the high temperature plasma, as seen in the initial 
phase of the M1 flare with the CDS on SOHO.


In this paper we focus primarily on the nature of the blue shifted
emission from high temperature plasma and its relationship with 
 the cool chromospheric emission and the hot flare emission.
In Sect.~2 we summarise the observing campaign for the 22nd October
2002.
In Sect.~3 we present the CDS observations and results in detail, 
paying particular attention to the measurement of wavelength shifts.
In Sect.~4 we discuss the  VTT H$\alpha$ observations and the method used
to derive chromospheric flows from the spectral line profiles. 
In Sect.~5 we draw
our conclusions from these observations.



\section{Co-ordinated observing campaign for 22nd October 2002}


\begin{figure}[!htbp]
\centerline{\epsfig{file=trace_1529.ps, width=8.5cm,angle=0 }}
%\centerline{\epsfig{file=trace_1529_ovl.ps, width=8.5cm,angle=0 }}
\caption{ TRACE 195 \AA\ image of the active region before the M1 flare.
%Bottom: the same but with contours ($\pm$ 200, 500, 1000 G) 
%from the SOHO/MDI magnetogram.
The fields of view of the TRACE sequence of images (Z1, see 
Fig.~\ref{fig:trace_seq}), of the region of the  M1 flare 
(Z2, see ~Fig.\ref{fig:trace_seq_zoom}),
and of the CDS raster are also shown.
}
  \label{fig:trace_over}
% Fig.~\ref{fig:trace_over}
\end{figure}


\begin{figure}[!htbp]
\centerline{\epsfig{file=trace_seq.ps, width=9.0cm , angle=0 }}
\caption{Sequence of TRACE 195~\AA\ difference images (negative) 
of the region Z1 (cf. Fig.~\ref{fig:trace_over}) 
and  centred during the peak of the M1 flare. 
The image that has been subtracted was taken at 15:29:59 UT,
before the start of the flare.
The timings indicate the initial exposure time 
for each TRACE image, and the corresponding nearest CDS exposure.
The location of the CDS slit is also shown.
}
  \label{fig:trace_seq}
% Fig.~\ref{fig:trace_seq}
\end{figure}




\begin{figure}[!htbp]
\centerline{\epsfig{file=trace_zoom.ps, width=6.5cm , angle=0 }}
\centerline{\epsfig{file=mdimap_zoom3.ps, width=6.cm , angle=0 }}
\caption{Top: sequence of three (T1,T2,T3) TRACE 195~\AA\ difference 
images (negative) in the region Z2 (cf. Fig.~\ref{fig:trace_over})  
of the  M1 flare. 
Bottom: magnetogram from SOHO/MDI. 
The reference pre-flare image was taken at 15:29:59 UT.
%Contours ($\pm$90 G) from the SOHO/MDI magnetograms are superimposed.
The areas P1, P2, P3 
and their surrounding mixed-polarity regions experienced changes 
during the flare. 
}
  \label{fig:trace_seq_zoom}
% Fig.~\ref{fig:trace_seq_zoom}
\end{figure}



The active region NOAA 0162 was mainly bipolar, quite extended, and
relatively active.
It was observed on October 22, 2002 with SOHO, TRACE, RHESSI,
and from the ground with  THEMIS and VTT.
The SOHO/CDS sequence of observations 
was extensive and ran during 08:02-13:07~UT and 15:18-18:13~UT
(see Table~\ref{tab:list_files}).
The CDS observations, described in detail below,
 consisted of 'rasters' during which the slit
was scanning the active region.

During October 22  the active region  produced at least 
3 flares, all of which were recorded by CDS.
A relatively strong  flare was recorded
during the CDS raster at 12:43--13:07 UT.
%
The second, main M1 flare is the focus of this paper. 
It occurred during the CDS raster at 
15:18:16--15:43:16 UT.
%
Fig.~\ref{fig:goes_plot} shows the X-ray light curves from GOES~8 
during the CDS scan.
The time derivative of the 
0.5--3 \AA\ light curve can be used as a proxy for the 
hard X-ray emission, and indicates a peak at 15:33 UT.
The timings  of some  CDS  slit positions and 
some TRACE and ground-based observations are also shown.


\begin{table}[ht]
\caption[]{CDS rasters analysed - 22 Oct 2002.}
\begin{center}
\begin{tabular}{ccl}
\hline\noalign{\smallskip}
 Time (UT) & CDS files  & CDS study \\
\hline\noalign{\smallskip}


8:02:50--8:27:54 & s26053r00  & SIG\_FT44 \\
 & s26054r00  & ARDIAG\_2 \\
 & s26054r01  & ARDIAG\_2 \\
 & s26055r00  & ARDIAG\_2 \\
 & s26056r00  & SIG\_FT44 \\

15:18:16--15:43:16  & s26058r00  & SIG\_FT44 \\
15:43:30--17:07:33  & s26059r00  & ARDIAG\_2 \\
17:08:03--18:32:46  & s26059r01  & ARDIAG\_2 \\

\hline\noalign{\smallskip}

\end{tabular}
\end{center}
\label{tab:list_files}
\end{table}
%Table~\ref{tab:list_files}

A third flare, class C2, was recorded later by another
CDS  raster during 15:43:30--17:07:33 UT,
and has been discussed in \cite{berlicki_etal:04}.
%
%THEMIS observed continuously between 15:41--17:26~UT 
%in the Na I ${\rm{D_{1}}}$ spectral line. 
%These observations were used to obtain  
% the line--of--sight (LOS) magnetic field  from the observed Stokes profiles.
%
VTT/MSDP observations in H$\alpha$ consisted of 
9 raster scans between 15:30 and 16:17 UT and a field-of-view
of about 380''$\times$170'' (exposure time of 0.5s).
The rastering was performed from east to west, in the 
opposite direction compared to CDS.
Below we will concentrate on the first VTT spectral raster,
which scanned the active region during 15:30-15:37 UT, i.e. 
in the impulsive phase. Results from the follow-up raster 
(gradual phase) have been presented by Berlicky et al. (2005).


TRACE 195~\AA\ images were obtained at high-cadence (40~s) 
between 15:29 and 15:41 UT (with a FOV
384''$\times$384'' and a 1'' resolution).
Before and after this time the TRACE observations were
obtained with slightly lower time resolution (50 s) and larger FOV
(512''$\times$512''). 
Before the M1 flare, the active region presented the usual 
characteristics, as shown by the TRACE 195~\AA\ images
(see  Fig.~\ref{fig:trace_over} for the full TRACE FOV).
That is the active region consisted of some unresolved emission 
and a large number of loops.

%
 Fig.~\ref{fig:trace_seq} shows a sequence of 
TRACE 195 {\AA} difference images in the central part of the 
TRACE FOV, where the flare took place.
The location and timings of the CDS slit are also indicated 
in  Fig.~\ref{fig:trace_seq}.
The first obvious signature  in the  TRACE 195 {\AA}
occurs  at 15:32:18 UT, somewhat later than the 
 increase in  the GOES X-ray flux.
However, a small brightening already appears at the same 
location on the 15:31:31 UT TRACE image.
The position of the flare and its relation to the underlying 
photospheric  magnetic field is shown in Fig.~\ref{fig:trace_seq_zoom},
where a zoom-up of the  same region is presented.


%
By 15:33 UT, the entire ribbon structures (located
on either side of the neutral line) brighten up in the TRACE 195~\AA\
image. The H$\alpha$ ribbons were already 
visible at 15:31 UT (see \citealt{berlicki_etal:04}). 
The two emissions appear co-spatial. 
It is possible that 
the TRACE 195~\AA\ emission in the brightest region consisted of 
 a system of low-lying loops that connected the two polarities.
%
% The East-side ribbon was compact and its position was almost
%constant during the flare while the West-side ribbon was 
%elongated and exhibited
%fast propagation along its length, from the emerging bipole towards the
%leading sunspot.
At the same time,
a fainter system of loops forms and expands in height, 
as shown in Fig.~\ref{fig:trace_seq}.
These loops are 
 rooted in the ribbons at specific points in the 
two opposite polarities.
%
The interpretation of the structures seen in 
the  TRACE 195~\AA\ images is  non-trivial, since
the 195~\AA\ band is mainly formed by the superposition of
continuum emission and line emission from 
 Fe~VIII, Fe~XII and Fe~XXIV, all of which are formed 
at very different temperatures.

CDS had the slit positioned over the main area of the flare 
during the peak and gradual phase (cf.  Fig.~\ref{fig:goes_plot}
and Fig.~\ref{fig:trace_seq}) and can therefore be used to 
estimate the temperatures of the various features recorded by TRACE.
%(note that these features are lasting long enough for this 
As shown below, CDS observed strong Fe~XIX emission in the 
locations of both the ribbons and the higher-lying loops
observed by TRACE, which suggests that  most likely 
the TRACE images are dominated by hot emission from Fe~XXIV.


SOHO/MDI took high-cadence (1m)
 full-disk magnetograms and dopplergrams during the 
observations of this active region.
The flare was located above a region of strong mixed polarity,
in  particular in the middle of the unipolar 
areas P1, P2, P3 indicated in Fig.~\ref{fig:trace_seq_zoom}.
The magnetograms do not show any dramatic changes during
the flare.
 However, the three unipolar regions did experience measurable
changes at the same time and positions in the flare as seen by TRACE.
 The magnetic fragment P1 (of the leading 
polarity), embedded in a region of opposite polarity, 
is the first one to show (at 15:31 UT) a drop of 100 gauss
The entire surrounding area subsequently develops changes (indicating 
some flux cancellation). 
Unfortunately most variations appear to occur on  small 
spatial scales and are at the limit of the sensitivity of the 
instrument (20-30 gauss). 
Furthermore, the measured values are time-averaged line-of-sight 
components, and it is therefore difficult to attempt a detailed 
reconstruction of the topology of the magnetic field and its variations.


%Note, however, that much higher magnetic fluxes are present in the 
%active region, and none of them (with the exception of the 
%large sunspot) show any signs of variations.
%Further, the fact that the variations are co-temporal 
%(within the 1m resolution) and in the same area as in the 
%TRACE images cannot be coincidental.


\section{CDS instrument and data analysis}

The CDS observations were performed with the 
normal incidence spectrometer (NIS).
%The slits are  normally oriented along the solar N-S direction.
The internal scan mirror allows CDS  to scan an 
 area of 4\arcmin\ square within a single `raster', from west to east. 
The N-S coverage was provided by the use of 
the long slit 4\arcsec x240\arcsec. 

The NIS is composed of two gratings, that  
produce two wave-bands  
(NIS~1: 308  -- 379 \AA\ and NIS~2: 513 -- 633 \AA;).
Telemetry constraints made it necessary to
pre-select a number of NIS wavelength windows to be extracted
on-board  from the spectra, before the data are 
telemetered to the ground. 


The CDS observing sequences were specifically designed 
to study the spatial and spectral characteristics of small flares.
They included a range of spectral lines from different ions and in
particular the high temperature flare line, Fe~XIX 592.2~\AA.
%
CDS observed NOAA 0162 
with the two types of rasters (see Table~\ref{tab:list_files}):
1) a slow diagnostic raster
(1hour 20 m, CDS study name ARDIAG\_2)
that extracts many diagnostic lines;
 2) a fast raster
(slit exposure of 18s for a total duration of 
20 m, CDS study name SIG\_FT44) that 
only extracts a few lines, selected to cover a wide  range of
temperatures (see Table~\ref{tab:list_lines}).





\begin{table*}[ht]
\caption[]{List of the main lines used for the wavelength calibration.
The first row gives the reference wavelengths ($\lambda_{st}$, \AA) adopted.
The second row gives the wavelengths ($\lambda_{1}$, \AA) as obtained from the 
pre-flare raster s26055 and the calibration curve adopted.
The third row gives the corresponding relative velocities ($v_1$, km/s).
The following two rows give wavelengths ($\lambda_{2}$, \AA) and
relative velocities ($v_2$, km/s) as obtained 
from the post-flare raster s26059 and using the same 
  calibration curve. Note the small differences in the 
measured wavelengths, giving confidence in the stability of the 
wavelength calibration.
}
\begin{center}
\footnotesize
%\tiny
%\scriptsize
\begin{tabular}{@{}lcccccccccccc@{}}
%[511.455, 0.116225, 7.71135e-07]
\hline\noalign{\smallskip}
\hline\noalign{\smallskip}
                     & Si XII & He I   & Ne IV  & Al XI & O IV & O IV&
Ca X     & Ne VI   & Ne VI & Fe XII & Mg X & O V \\

$\lambda_{st}$  & 520.665& 522.210&  542.07& 550.06& 553.34& 555.263&
 557.765& 558.594& 562.803& 592.60& 624.941& 629.732 \\

$\lambda_1$ & 520.675 & 522.208   &    542.031  &     550.070   &    553.365  &     555.252    &   557.758     &  558.590    &   562.791    &   592.590   &    624.936
   &    629.735 \\

$v_1$  & 5 &   -1 &  -21  &   5 &   13 &   -6 &   -3  &  -2 &   -6 &   -5 &   -2  &   1 \\

$\lambda_2$  & 520.664 & 522.211 & 542.061 & 550.066 &
 553.354 & 555.260 & 557.769 & 558.622 & 562.812 & 592.597 & 624.932 & 629.747 \\

$v_2$      & 0 &   0 &  -4 &   3  &  7 &  -1  &  1 &  15  &  5  & -1 & -4  &  6 \\

\hline\noalign{\smallskip}
\end{tabular}
\end{center}
\normalsize
\label{tab:wvl_cal}
\end{table*}
%Table~\ref{tab:wvl_cal}



\begin{table}
\caption{Some of the prominent  lines in the CDS/NIS wavelength bands,
used in this paper.
Wavelengths are taken from the CHIANTI database unless otherwise
stated.
Those indicated with an asterisk have been used for the 
wavelength calibration.
When only one decimal 
place is indicated, the lines are blends of at least two transitions.
The approximate temperature of maximum ion abundance 
(in ionization equilibrium) is also
indicated, together with the channel and the CDS study that 
recorded the lines (FR: fast raster; DR: diagnostic raster, see text).
%
%*************************
}
\begin{center}
%\footnotesize
%\tiny
%\scriptsize

\begin{tabular}{llrccc}
\hline\noalign{\smallskip}
Ion & wavelength (\AA) & log T$_e$  & NIS  &  \\
\hline\noalign{\smallskip}

Fe XIV & 334.172 & 6.2 &   1  &  DR \\

Si X & 347.402 & 6.1 &   1  &  FR,DR \\
Si X & 356.0 & 6.1 &   1  &  FR,DR \\

Fe XVI & 360.761 & 6.3 &   1  &  DR \\
Fe XII & 364.467  & 6.1 &   1  &  DR \\

Mg VII  & 367.7 & 5.8 &   1  &  DR \\
Mg IX  & 368.070 & 6.0 &   1  & DR  \\

Si XII & 520.665$^*$ & 6.3 &   2  & FR,DR \\
He I  & 522.210$^*$ & 4.5 &   2  & FR,DR \\

Ne IV & 542.07$^*$ & 5.3 &   2  & DR \\

Al XI & 550.06$^*$ & 6.2 &   2  & DR \\

O IV  & 553.34$^*$ & 5.3 &   2  & FR,DR \\
O IV  & 554.076 & 5.3 &   2  & FR,DR \\
O IV  & 554.513 & 5.3 &   2  & FR,DR \\
O IV & 555.263$^*$ & 5.3 &   2  & FR,DR \\

Ca X & 557.765$^*$ & 5.8 &   2  &  DR\\
Ne VI  & 558.594$^*$ & 5.6 &   2  &  DR\\

Ne VII  & 561.728 & 5.7 &   2  &  DR\\
Ne VI  & 562.803$^*$ & 5.6 &   2  &  DR \\

Fe XIX & 592.16 & 6.9 &  2  & FR,DR \\

Fe XII & 592.60$^*$ & 6.1 &  2  & FR,DR \\

Mg X & 624.941$^*$ & 6.0 &   2  & FR,DR \\
O IV &  625.853   & 5.3 &   2  & FR,DR \\
O V  & 629.732$^*$ & 5.4 &   2  & FR,DR \\
\hline\noalign{\smallskip}
\end{tabular}
\end{center}
\normalsize
\label{tab:list_lines}
\end{table}
% Table~\ref{tab:list_lines}



One important aspect of CDS is that lines emitted from 
the chromosphere up to the corona are observed. 
This is very useful for a proper co-alignment 
with other datasets, e.g. photospheric magnetic fields 
from SOHO/MDI, and coronal images from TRACE.
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 and high count rates, cosmic ray removal.
%
Multiple-line-fitting techniques have been used
to determine the moments of each of the spectral 
lines in the spectra by removing the {\it `continuum' }
 intensity, which is mainly  due to scattered light in the NIS.
%
In SOHO post-recovery spectra,  the CDS/NIS line profiles
are no longer gaussians, and show a broadened and asymmetrical
long wavelength wing, which is approximately
constant and can therefore be accounted for.
%
The CDS calibration of \cite{delzanna01_cdscal}
has been adopted here,  with a correction to take
into account the effects on the NIS detector of the gain depression caused by
exposures with the wide slit.




\subsection{Wavelength calibration}


For a correct interpretation of the velocities observed,
an accurate wavelength calibration is necessary.
%
CDS does not have an absolute wavelength calibration.
%
Changes in the  wavelength calibration 
have been observed (up to 0.1~\AA\ in NIS 2), 
and caused primarily by on-board
temperature changes (and their history),
which are due to the variable illumination
received from the Sun (D. Pike, priv. comm.).
Also, the wavelength calibration often 
requires a `scan mirror correction', 
which also  appears to vary.
%
In addition, there are no photospheric or chromospheric lines within 
the CDS wavelength bands (except He~I).

The above means that the calibration has to be done using the 
observed transition region and coronal lines, and that 
it is necessary to obtain/check the
wavelength calibration for each set of observations.
We note, however, that in the literature it is 
standard practice (see, e.g. \citealt{brosius_phillips:04})
 to only make relative wavelength measurements, 
by using as reference values those obtained within 
the same set of observations, and without performing a 
wavelength calibration.

For the calibration, it is important to have accurate `standard'
rest wavelengths for the spectral lines concerned, which is 
a non-trivial issue if very accurate measurements are 
sought. 
We have adopted the values shown in Table~\ref{tab:list_lines}.
For most lines we adopt the wavelengths as in the CHIANTI 
database, which in turn has values mostly from NIST. 
We note, however, that 
the values in NIST are taken from the literature, where 
many measurements have been obtained from solar spectra.
A simple survey in the literature shows that for most 
spectral lines wavelength measurements are inconsistent
(i.e. differences are larger than the uncertainties in the 
measurements). 
These issues are discussed in a series of papers 
where atomic data (including wavelengths) are 
benchmarked on an ion-by-ion basis (cf.
\citealt{del_zanna_etal_fex})


The case of the important \ion{Fe}{xix} {592.2}\AA\ line
is emblematic. A large number of different measurements can 
be found in the literature.
% 
%This line was observed by ...
%and identified for the first time by ... 
%However, accurate measurements are sparse in the literature,
%the few available being inconsistent with each other.
%
Our calibration, based on a small sample,
provides a wavelength of 592.16 \AA, in excellent 
agreement with the accurate measurements of \cite{widing:78}, 
but different from e.g. the value previously reported 
by  \cite{sandlin_etal:76},  592.24~\AA.
We note that \cite{widing:78} averaged many flare observations taken with 
the NRL spectroheliogram on Skylab.
On the other hand, 
\cite{lawson_peacock:80}, in their review of the $n=2 \to n=2$ 
transitions, based on  solar and
laboratory measurements,  suggested an energy 
for the \orb[2s 2] \orb[2p 4] $^1$D$_{2}$ level 
of 168770 \cm. This energy is equivalent to a rest wavelength of 592.522 \AA,
i.e. at  the considerable difference of  
 183 km/s, compared to our result, 592.16~\AA.
%
% {sandlin_etal:76}  592.24 
%
An extra complication is due to the presence, in the 
red wing of the  \ion{Fe}{xix} line, of a coronal 
line only recently identified as due to \ion{Fe}{xii}
(see Fig.~\ref{fig:fe_xix_win} below and 
\citealt{delzanna02_cospar}, \citealt{del_zanna_mason_fexii}).


To calibrate the CDS observations, 
average CDS spectra containing a large number of lines 
(the diagnostic rasters) of two rasters taken before and 
after the flare have been analysed and wavelength-calibrated.
We assumed for the CDS pixel-wavelength calibration 
a quadratic function.
The  wavelength calibration was remarkably stable over the 
period. Table~\ref{tab:wvl_cal} shows that in most cases the 
line centres were within a few km/s the assumed reference 
values.

In terms of uncertainties, we note that  
the positions of the line centres can be calculated with 
approximately 1/10 of a pixel accuracy 
(1  pixel in the NIS~2 channel is approximately equivalent to 0.11 \AA),
and the average differences between the predicted and 
standard wavelengths are of the same order, 0.01 \AA.
This accuracy results in an uncertainty 
in the derivation of velocities of the order of 
10 km/s.
An additional systematic uncertainty is due to the 
spatially-dependent effects due to the geometrical 
distortions within the instrument. 
This is more difficult to quantify, because it depends on the 
type of relative measurement that is done, but is again 
of the same order.
Finally, we note that the resolution of the CDS instrument 
does not allow simultaneous measurements of 
stationary and non-stationary components. Hence, our 
results refer in most cases to the average main displacement of 
each line profile. 
This is not true in the case of the Fe~XIX emission, since 
it is present only during flares.
Signicant line broadening can however be 
measured even in post-recovery spectra.

\begin{figure}[!htbp]
\centerline{
\epsfig{file=s26058r00_HeI_int_mdi.ps, width=4.7cm,angle=0 }
%\epsfig{file=s26058r00_HeI_pos_mdi.ps, width=8.0cm,angle=0}
\epsfig{file=s26058r00_OV_int_mdi.ps, width=4.7cm,angle=0}}
%
\centerline{\epsfig{file=s26058r00_SiXII_int_mdi.ps, width=4.7cm,angle=0 }
\epsfig{file=s26058r00_FeXIX_int_mdi.ps, width=4.7cm,angle=0 }}
\caption{Monochromatic images obtained from the CDS raster.
Contours from the SOHO/MDI magnetograms are superimposed.
}
  \label{fig:cds_ima}
% Fig.~\ref{fig:cds_ima}
\end{figure}


\begin{figure}[!htbp]
\centerline{\epsfig{file=trace_195_1539_FeXIX_int.ps, width=6.0cm,angle=90 }}
\caption{ TRACE 195 \AA\ difference (negative) image
 with contours of Fe XIX intensity.
}
  \label{fig:trace_fe_xix_int}
% Fig.~\ref{fig:trace_fe_xix_int}
\end{figure}

\begin{figure}[!htbp]
%\centerline{\epsfig{file=s26058r00_SiXII_pos_mdi.ps, width=8.0cm,angle=0}}
%\centerline{\epsfig{file=s26058r00_SiXII_pos.ps, width=8.0cm,angle=0}}
\centerline{\epsfig{file=s26058r00_sixii_areas.ps, width=8.0cm,angle=0}}
%\centerline{\epsfig{file=s26058r00_FeXIX_int_SiXII_pos.ps, width=8.0cm,angle=0}}
\caption{\sixii\  velocity map in the flare region,
 with the areas selected for further  analysis.
Areas A,B,C correspond to regions of strong blue-shifts,
while area D of red-shift. Region LP corresponds to 
the area of peak \fexix\ intensity, while region L1 is where 
the central part of the flare loop is located. 
Region BG is a `background' reference region.}
  \label{fig:si_xii_pos}
% Fig.~\ref{fig:si_xii_pos}

\end{figure}


%\begin{figure}[!htbp]
%\centerline{\epsfig{file=s26058r00_FeXIX_pos_SiXII_pos.ps, width=7.0cm,angle=0}}
%\centerline{\epsfig{file=trace_195_154039_SiXII_pos.ps, width=8.0cm,angle=0}}
%\centerline{\epsfig{file=s26058r00_FeXIX_int_SiXII_pos.ps, width=8.0cm,angle=0}}
%\caption{ Fe~XIX image with superimposed the
%  contours of the Si~XII blue- and red-shifts. }
%  \label{fig:over_SiXII_pos}
%% Fig.~\ref{fig:over_SiXII_pos}
%\end{figure}


\subsection{Overview of the CDS results}


\begin{figure}[!htbp]
%\centerline{\epsfig{file=s26058r00_oiv_win.ps, width=5.cm , angle=90 }}
\centerline{\epsfig{file=s26058r00_mgx_win.ps, width=5.5cm , angle=90 }}
\caption{Spectra of three different regions, A,B,BG. 
Note that in region A the profiles of the strongest lines are 
nearly saturated. The O~IV line has been used for density measurements
at transition-region temperatures.}
  \label{fig:oiv_win}
\end{figure}
% Fig.~\ref{fig:oiv_win}

\begin{figure}[!htbp]
\centerline{\epsfig{file=s26058r00_fe_xix_win.ps, width=5.5cm,    angle=90 }}
\caption{Spectra of the region L1 (peak flare),
and of the background BG (multiplied by ten), showing the
\fexix\ and the previously unidentified Fe~XII line.
Outside of the flaring region, all the intensity is due to Fe~XII.}
  \label{fig:fe_xix_win}
\end{figure}
% Fig.~\ref{fig:fe_xix_win}

\begin{figure}[!htbp]
%\centerline{\epsfig{file=s26058r00_he_i_prof.ps, width=5.5cm , angle=90 }}
\centerline{\epsfig{file=s26058r00_si_xii_prof.ps, width=5.cm , angle=90 }}
\centerline{\epsfig{file=s26058r00_fe_xix_prof.ps, width=5.cm , angle=90 }}
\caption{Line profiles in regions A,B,D showing the evident
 blue-shifts in the coronal (Si~XII) and flare (Fe~XIX) lines in
 regions A,B. }
  \label{fig:prof}
\end{figure}
% Fig.~\ref{fig:prof}




Fig.~\ref{fig:cds_ima} shows monochromatic images obtained from the
CDS raster, in lines emitted over a broad temperature range.
The ribbon structure is so bright in He~I and also at 
transition region temperatures (e.g. O~V) that the brightest 
lines saturated the detector in some positions 
(region A, defined below). 
The hot emission from Fe~XIX is only present in the 
ribbon and high-lying loop  structures seen in the TRACE 195 \AA\ band 
(see Fig.~\ref{fig:trace_fe_xix_int}).

%This suggests that the brightest emission seen in the 
%TRACE 195 \AA\ band is dominated by Fe~XXIV emission.

The velocity maps show a consistent pattern.
Large velocities, mainly upflows, are only visible in 
lines formed above 1~MK, and in the  Fe~XIX flare line.
Fig.~\ref{fig:si_xii_pos} (top) shows as an example a velocity
map in the flare area obtained from a Si~XII line,
and the regions selected for further study. 
Line fitting was performed on the 
 averaged spectra of these selected regions.
Figs.~\ref{fig:oiv_win},\ref{fig:fe_xix_win},\ref{fig:prof} 
show some line profiles in a few selected regions.
Table~\ref{tab:areas} presents some results of the line 
fitting, in particular  velocities,  intensities and 
line widths.

%
Regions A,B were  associated with 
the brightest emission in the chromospheric-TR lines,
and with strong upflows in the Fe~XIX flare line.
Region A was observed by CDS during 15:33:20--15:34:10 UT  
(see Fig.~\ref{fig:trace_seq} for the correspondence with TRACE),
 i.e. in the impulsive phase of the flare
 (Fig.~\ref{fig:goes_plot}).
%
Region B was observed during 
15:35:00--15:35:50 UT, i.e. at the time of the peak
X-ray  emission (Fig.~\ref{fig:goes_plot}). 

As observed in other cases (\citealt{del_zanna_etal_prague}), upflows are also 
seen in the lower-temperature Si~XII and Mg~X lines. 
Our findings are contrary to those of 
\cite{brosius_phillips:04}, that  
discouraged the use of  Mg~X lines because of 
blends with O~IV emission. 

In terms of line widths, aside from a few notable exceptions,
the  full-width at half maximum ($FWHM$) values are close to 
the value of the instrumental width (0.5~\AA\ in post-recovery 
spectra). The exceptions are in regions A,B,D and mostly in 
the coronal lines. Interestingly, the Fe~XIX line has a large
width in all of the regions, suggestive of large turbulent 
motions.
As also noted by \cite{brosius_phillips:04}, the 
observed blue-shifts are not superimposed on a stationary 
component, but appear as 
shifts of the entire line profiles
 in Mg~X, Si~XII and Fe~XIX  (see Fig.\ref{fig:prof}). 


We also find upflows  (although smaller) in the TR lines
in region B.
Interestingly, in region A, during the impulsive phase, He~I, O~IV and O~V
 lines show small downflows. However, these results should 
be considered with some caution because the detectors were nearly 
saturated (completely saturated in the O~V case).


The other regions were close to the flare site.
Region L1 shows near-zero line-of-sight velocities 
in all lines.
The region LP corresponds to the area of peak X-ray  emission,
as judged from the TRACE 195~\AA\ images and the 
Fe~XIX emission, observed by CDS 7 minutes later, around 15:42 UT.
It is interesting to notice that, on the opposite sides of the 
peak emission, flows of opposite directions are observed 
(region C, upflow; region D, downflow) in all the lines, emitted 
from TR to flare temperatures.
%

Region BG was chosen as a background reference 
region, not interested by the flare.
The lines 
show near-zero line-of-sight velocities.  
Note that the intensity of the Fe~XIX line is 
an upper limit (the emission being dominated by the  Fe~XII).
We have compared the line intensities of the background 
region with other observations of the quiet Sun, and found that,
for the transition region lines, the values are very close 
to those of the average quiet Sun. 



\begin{table*}[!htbp]
\caption[]{Results from the averaged spectra of the various 
regions displayed in Fig.~\ref{fig:si_xii_pos}.
For each of the selected spectral lines, we display
the averaged velocities  $v$ (in km/s; negative corresponds to blue-shift),
the calibrated line intensities $I$ (erg cm$^{-2}$ s$^{-1}$
sr$^{-1}$), and the full-width at half maximum  $FWHM$ (\AA).
O IV(a) indicates the 553.4~\AA\ line, while 
O IV(b) indicates the 625.8~\AA\ line, used to obtain the 
electron densities $N_{\rm e}$ (cm$^{-3}$).
Large blue-shifts are mostly observed in regions A,B,C.
Note, however, that some of the lines are close to saturation
in region A, and values are uncertain (they are indicated in 
parentheses).
 }
\begin{center}
%\footnotesize
%\tiny
%\scriptsize

\begin{tabular}{rrrrrrrr}
\hline\noalign{\smallskip}

\hline\noalign{\smallskip}
Region       & A  &   B  & C   & D   & L1 & LP  & BG  \\
\hline\noalign{\smallskip}
%
 He I $v$  &  13  & -27  & -30 & -17   & -17   & -36   &  -22   \\
      $I$  & 2393 & 770  & 112 & 377   & 751   & 1293   &  16   \\
   $FWHM$  & 0.61 & 0.53 & 0.51& 0.56  & 0.51   & 0.51   &  0.50    \\
\\
% 553.4:
 O IV(a) $v$ & (33) & -21  & -24 & 13    &  0   &  -20  &  0    \\
      $I$  & (9818) & 598  & 80  & 237   & 252   & 686   & 33    \\
%   $FWHM$  & (0.65) & 0.64 & 0.65& 0.65  & 0.65   & 0.65   &  0.56    \\
%625.853 :
\\
 O IV(b) $I$ & 2333 & 111  &  -  &  -    &  -  & -   &   -   \\

$N_{\rm e}$ (O~IV)  & 2 10$^{11}$ & 0.9 10$^{11}$ & - & - & -  & -  & - \\
\\
 O V  $v$  & (23)   & -25  & -21 &  0    &  -12   & -28   &  -4      \\
      $I$  & (33226) & 5038 & 721 & 2490 & 2840   & 6782   & 387    \\
   $FWHM$  & (0.70) & 0.50 & 0.50& 0.55  & 0.50   & 0.50   &  0.50   \\
\\
Mg X  $v$  & (14)   & -62  & -51 & 29    & -17   &  -37  &  -5    \\
      $I$  & (5910) & 1630 & 656 & 1205  & 1781  & 17093   & 118    \\
   $FWHM$  & (0.80) & 0.59 & 0.53& 0.58  & 0.58  &  0.52  &   0.52 \\
\\
 Si XII $v$ & -64  & -84  & -54  &  17   &   -6  & -40   & -10      \\
      $I$  & 2756 & 1847 & 839   &  1378 & 2818  & 16748 & 63     \\
   $FWHM$  &  0.59& 0.55 & 0.50  & 0.55  & 0.51  & 0.50   & 0.50    \\
\\
 Fe XIX $v$ & -141 & -140 & -56  & 22    & -13  & -26    & -     \\  
       $I$  & 294  & 272  & 694  & 1271  & 2721 &  7544  & $\le$ 0.5     \\
   $FWHM$   & 0.68 & 0.65 & 0.56 & 0.59  & 0.57 &  0.55  &  -   \\

%  &  & & & & & & & \\
 
%

\hline\noalign{\smallskip}
\end{tabular}
\end{center}
\normalsize
\label{tab:areas}
\end{table*}
% Table~\ref{tab:areas}

%
The O~IV lines (cf. Fig.~\ref{fig:oiv_win}) 
are an useful density diagnostic at 
transition region temperatures, and in the locations 
of strong brightenings. 
For both the upflow regions A and B we obtain  
 averaged line-of sight electron 
densities of $\simeq$10$^{11}$ cm$^{-3}$.
Interestingly,  the O~IV intensities increase  by 
factors of $\simeq$300, 20, while the density 
increases by an order of magnitude compared to 
quiet Sun values (cf. \citealt{delzanna_jgr99a}).
This means that the filling factor of the two regions 
could be different. This is an added complexity that 
a detailed flare interpretation should take into account.
%Accurate measurements of filling factor require, however,
%much higher spatial resolution than the CDS one.







\section{ H$\alpha$ line profiles}

For the analysis of the H$\alpha$ Dopplershifts we used the observations of 
VTT/MSDP obtained at 15:30-15:36 UT. We selected areas (Fig.~\ref{fig:vtt})
corresponding to the regions observed by SOHO/CDS 
 (see Fig.~\ref{fig:vtt_over}). 
Regions 6,7,8 in the H$\alpha$ spectra 
correspond to region A in the CDS spectra; 4,5 to region B;
and 3 to region L1.

The mean profile of H$\alpha$ line from each analysed area was obtained
by averaging the spectra from nine pixels contained
inside a square box $0.75 \times 0.75''$, the distance between
two neighbouring pixels being equal to $0.25''$. We used these profiles 
to derive the velocity structure of the flaring chromosphere in the 
eight respective bright kernels. Fig.~\ref{fig:vtt_prof} presents 
these averaged profiles (continuous lines). These profiles are typical 
for flares, which often have double-peaked emission with a central 
reversal. All observed profiles are asymmetrical, indicating the 
presence of vertical motions in the chromosphere. 



\begin{figure}[!htbp]
\centering
\includegraphics[width=9cm]{vtt_points_landscape_1530_times.ps}
\caption{VTT/MSDP image of AR 0162 taken between 15:30 and 15:36 UT with dots
 which mark the areas where the H$\alpha$ line profiles were obtained.
 The timing and scanning direction are also indicated.}
 \label{fig:vtt}
% Fig.~\ref{fig:vtt}
\end{figure}


\begin{figure}[!htbp]
%\centerline{\epsfig{file=cds1518_heI_vtt1530_dots.ps, width=6.5cm,angle=0}}
\centerline{\epsfig{file=cds1518_vel_si_vtt1530_dots.ps, width=6.5cm,angle=0}}
\caption{{ 
%SOHO/CDS negative image of AR 0162 made between 15:18 and 15:43 UT in 
%         the He I 522.2 {\AA} line (upper panel) with dots (1-8) which mark 
%      the   corresponding areas where the H$\alpha$ line profiles were obtained.
%         In the lower panel we present the
%
 SOHO/CDS velocity map in the  Si XII 520.7~{\AA} line obtained 
between 15:18 and 15:43 UT. 
The darker structures represent upflows.
The dots (1-8) mark the 
areas where the H$\alpha$ line profiles were obtained
 The contours represent the H$\alpha$ flare ribbons
         as observed by VTT during 15:30-15:36 UT.}}
 \label{fig:vtt_over}
% Fig.~\ref{fig:vtt_over}
\end{figure}


In order to construct the models of the solar chromosphere corresponding
to the analysed areas of the flare, we follow the method discussed
in \cite{berlicki_heinzel:04}, which consists of two steps.
First,  a grid of static models is calculated
with the non-LTE code developed by \cite{heinzel:95},
and modified for flare conditions (\citealt{berlicki_heinzel:04}),
 by varying two parameters $m_{\mathrm{0}}$ and $\Delta T$.
A least-square fit is then applied to the
 observed and synthetic H$\alpha$ line profiles.
The $m_{\mathrm{0}}$ and $\Delta T$ values obtained 
are given in \tab{modpar1} (second and third columns).


Second, to investigate the influence of the macroscopic velocity field
on the emergent H$\alpha$ line profiles we adopted the formula used by
\cite{mihalas:78} to describe an expanding atmosphere.
This formula was previously used in the modelling of the velocity field 
in flaring atmosphere by \cite{nejezchleba:98} and \cite{berlicki_etal:05}.
The vertical velocity as a function of the H$\alpha$ line-centre optical depth
was defined as:
\begin{equation}
V(\tau)=\frac{2V_{0}}{1+\tau/\tau_{\mathrm{m}}} .
\end{equation}
$\tau_{\mathrm{m}}$ corresponds to the height where the velocity is equal
to $V_{0}$. If $\tau \ll \tau_{\mathrm{m}}$, which
corresponds to upper layers of the chromosphere, the velocity
goes to $2V_{0}$, and for $\tau \gg \tau_{\mathrm{m}}$
(close to the photosphere) the velocity reaches zero. As was shown 
by \cite{berlicki_etal:04} such a vertical structure of the velocity 
field seems to be a reasonable representation of an evaporating atmosphere. 
A second grid of  synthetic H$\alpha$ line profiles was computed using different 
values of $V_{\mathrm{0}}$ and $\tau_{\mathrm{m}}$, and 
a best-fit with each observed profile performed.
The fitted  profiles are shown in Fig.~\ref{fig:vtt_prof}.
 The parameters
describing the velocity field in the various areas are  also given in
\tab{modpar1} (forth 
and fifth columns).
The vertical velocity parameter $V_{0}$ was corrected with the value of direction
cosine $\mu=0.8$ assuming that the velocity vector is perpendicular to 
the solar surface. 
Negative values of $V_{0}$ parameter means upward flows. 


\subsection{Models of the flaring chromosphere}


\begin{figure}[!htbp]
\centering
\centerline{\epsfig{file=22oct_prof_1530b.ps, width=7.5cm,angle=0}}
%\includegraphics[width=12cm,bb=105 190 467 764,clip=]{22oct_prof_1530b.ps} 
%\psfig{figure=22oct_prof_1530.ps,width=12cm}
\caption{The observed (continuous lines) and fitted (dashed lines)
         H$\alpha$ line profiles in areas No. 1 - 8  of the flare. With
         dotted line we also plotted the mean profile of the quiet-Sun
         chromosphere taken from David (1961) for $\mu=0.8$. For areas
         1 and 2 we did not show fitted profiles (see text).}
 \label{fig:vtt_prof}
% Fig.~\ref{fig:vtt_prof}
\end{figure}



\noindent
\begin{table}[!htbp]
\centering
\caption{The parameters $m_{0}$ and $\Delta T$ of the static models 
         and $\mathrm{log}\;\tau_{\mathrm{m}}$ and the vertical velocity 
         $V_{0}$ of the models with velocity obtained from the fitting of
         synthetic H$\alpha$ line profiles with observed ones. The negative 
         sign of the velocity means upflow.}
\vspace{0.5cm}
\begin{tabular}{*{5}{c}} \hline\hline 
 {$\rm{Area}$} & {$m_{0}{\rm [g\;cm^{-2}]}$} & {$\Delta T{\rm [K]}$} & {$\rm{log}\;\tau_{m}$} & {$V_{0} \rm{[km\;s^{-1}]}$} \\
\hline
1 & $1.0 \cdot 10^{-3}$ & $+100$ & $-0.4$ &  downflow (?) \\
2 & $1.0 \cdot 10^{-3}$ & $-200$ & $-0.2$ &  downflow (?) \\
3 & $1.0 \cdot 10^{-3}$ & $ +50$ & $+0.4$ & $+10.0$ \\
4 & $8.0 \cdot 10^{-4}$ & $-100$ & $+1.0$ & $ -1.9$ \\
5 & $8.5 \cdot 10^{-4}$ & $-200$ & $+0.1$ & $ -3.8$ \\
6 & $8.5 \cdot 10^{-4}$ & $-250$ & $-0.8$ & $ -9.4$ \\
7 & $9.5 \cdot 10^{-4}$ & $-300$ & $-0.1$ & $ -5.6$ \\
8 & $1.0 \cdot 10^{-3}$ & $-400$ & $ 0.0$ & $ -6.9$ \\ \hline

\end{tabular}
\label{modpar1}
\end{table}
\indent


Most of the profiles in the flare ribbon areas
exhibit a red asymmetry manifested by the slight increase of the intensity 
in the red wing around \mbox{$0.4-0.5$ {\AA}}.
In the case of self-reversed profiles the red asymmetry means an upflow. 
Among the eight analysed line profiles, one of them indicated downflow 
velocity and in five line profiles we found upward flows (see \tab{modpar}).
Interestingly, the same analysis performed on the following 
raster by \cite{berlicki_etal:05} indicates very similar upflows,
which suggests that they were a long-lasting feature.

Fig.~\ref{fig:vel} shows the  blue-shifts in region B 
observed by CDS, together with the corresponding value 
obtained from the H$\alpha$ measuerements.

\begin{figure}[!htbp]
\centerline{\epsfig{file=vel_temp.ps, width=8.5cm , angle=0 }}
\caption{Observed blue-shifts in region B}
  \label{fig:vel}
\end{figure}
% Fig.~\ref{fig:vel}


The profiles in areas 1 and 2 are very flat or asymmetric and cannot be 
fitted with satisfactory quality. The shape is consistent with  strong downflows.



The non-LTE codes allows us also to calculate
the electron densities at different heights of our
model atmospheres. The models of solar atmosphere
 contain the layers from the photosphere 
to the chromosphere-corona transition region (CCTR). The CCTR temperatures obtained
for the analysed areas 1-8 are $\simeq 10^{5}$ K. The electron
densities $N_{e}$ are $\simeq 1.0 \cdot 10^{12}\; \mathrm{cm^{-3}}$.

%These values are somewhat higher than the averaged densities obtained from the CDS lines
%(see Tab.4). However, one should remember that the O~IV lines
%are formed at temperatures somewhat higher than  H$\alpha$
%and therefore higher in the solar atmosphere, where the $N_{e}$ is likely to be lower.



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


Some interesting aspects of the flare have been 
observed and deserve further comment.
%

The  co-spatiality between the  CDS  Fe~XIX and the 
TRACE 195~\AA\ emissions suggests that 
the bulk of the  TRACE 195~\AA\ emission 
 at the flare site is due to Fe~XXIV.
The TRACE 195~\AA\ images, together with the 
SOHO/MDI magnetograms,  clearly indicate that 
the location of the peak X-ray emission recorded by 
GOES is above a  region of strong mixed polarity.
Imaging in a flare line at high-time resolution 
($<$ 1m) is obviously a requirement for future 
observations.

Interestingly, in the same region the MDI magnetograms show 
small (at the limit of the sensitivity) but significant
changes, co-temporal with the TRACE/GOES light curves.
The changes occur on short time-scales (minutes) and 
indicate permanent flux cancellation in the area
of highly-mixed polarities.
%
There is a large body of literature 
on ground-based and space magnetic field observations
before/during/after a flare. The results are 
controversial in the sense that changes 
in the magnetic fields have only been observed in a few cases.
Small variations based on MDI magnetograms,
but on longer timescales and on much larger 
flares have been reported (cf. \citealt{kosovichev_zharkova:01}).
It is therefore quite surprising to see measurable 
changes on timescales of one minute for a relatively 
small flare. 

However, for a global interpretation of the event,
one should keep in mind that the MDI measurements
only provide an estimate of the line-of-sight component
of the magnetic field.
Furthermore, some doubts have been expressed about 
the MDI measurements themselves. For example, 
accelerated particles and radiation can alter the continuum 
and the profile of the lines used to infer 
magnetic fields, with the effect that the standard 
methods can provide incorrect values for the 
magnetic fields (see, e.g. 
\citealt{ding_etal:02} for a discussion on MDI 
magnetograms). A correct interpretation would require 
full line profile and continuum measurements and
a non-LTE model which would take into account of all 
processes that can affect the formation of the observed 
line from which the magnetograms are obtained.
Future instruments such as those on SOLAR-B will 
have a much higher sensitivity and will
produce vector magnetographs, which will greatly 
enhance our understanding.

CDS has measured flows in this flare in spectral lines formed
over a range of temperatures (log T=5.4-6.9)
with an accuracy $\simeq$10 km/s. 
As seen in our observations of many other small flares, 
we find upflows in `hot' lines that are 
concentrated in small spatial areas, at the footpoints 
of coronal/flare loops.
Stronger upflows are in the Fe~XIX line,
indicating that (as we already knew) most of the action 
takes place in the hotter lines. 
The large broadenings of the Fe~XIX line
are in fact indicative of the presence of turbulent motions,
in addition to the overall flows.


Coronal images alone give no indications whatsoever on where 
flows could be located. The strongest flows seem to be observed 
in regions of low emission. 
Spectral observations of flare lines at high-time 
resolutions are needed.  SOLAR-B EIS will have 
such capabilities.


%This evidence counts in favour of fast raster scans over an 
%active region, instead of a sit-and-stare approach.
There seems to be no clear one-to-one correspondence between the
TR brightenings and the flows observed in the coronal emission.
Fast imaging in TR lines would be useful, considering that 
most of the radiative losses occur in the transition region.
Spectral observations with high resolutions would be needed to 
measure filling factors. 
Unfortunately, no planned instrument will have such capabilities.

An important diagnostic tool available with spectroscopic instruments 
is the possibility
of measuring electron densities over a range of temperatures.
 We were able to use the  CDS
O~IV diagnostic lines to obtain electron densities of 
the order of 10$^{11}$ cm$^{-3}$ for the regions A,B.
These are somewhat lower that the values of 10$^{12}$ cm$^{-3}$
derived from the H$\alpha$ data.


The Doppler-shifts of the chromospheric ribbons (10$^4$ K )
have been derived from the profiles of H$\alpha$ using a non-LTE
radiative transfer code with a flare model
and a velocity function (\citealt{berlicki_etal:05}). 
%
To our knowledge, this is only the second time (after 
 \citealt{teriaca_etal:03}) that simultaneous and co-spatial
 chromospheric and coronal spectroscopic observations showing
line shifts during flare are reported.
In the C2 eruptive flare studied by \cite{teriaca_etal:03}
upflows in coronal/flare lines corresponded 
to  downflows in the chromosphere as seen in the 
He~I 10830~\AA\ line.
%
In our case we have a different picture.



The enhanced intensity of the red wing observed for the H$\alpha$ line profiles in 
areas 1 and 2 suggests that in these areas we observe large redshifts. 
These H$\alpha$ line profiles were observed at 15:31 UT. This time 
corresponds to the increase of the GOES soft X-ray flux, which reached the maximum  
at around 15:35 UT. This suggests that these H$\alpha$ profiles were observed 
during the impulsive phase. 
Such a type of asymmetry of H$\alpha$ line profiles was reported in \cite{ichimoto_kurokawa:84} 
and it can be caused by downward motion of the flaring chromosphere with 
velocities more than $40\;\mathrm{km\;s^{-1}}$. \cite{fisher_etal:85c} suggests 
that such a downflow of cool chromospheric flaring plasma (chromospheric condensations) 
is caused by the enhanced pressure produced in the region heated by non-thermal electrons 
accelerated in the corona during the impulsive phase of solar flares. 
The downflow observed in cool chromospheric H$\alpha$ line can be a reaction on 
explosive evaporation of hot plasma frequently observed in X-ray lines.


The areas that were observed simultaneously 
by CDS (A,B) and in H$\alpha$ (regions 5,6,7,8) during the peak 
X-ray emission presented blueshifts which could
correspond to the `gentle evaporation' of the chromospheric plasma.
The blueshifts are of the order of 5 -- 10 km/s in H$\alpha$, 
27 km/s in He I, 
60 -- 88 km/s in Mg~X, Si~XII, and 140 km/s in Fe XIX in 
region B. Region A shows instead downflows in the cooler lines
observed by CDS, but the velocities are questionable 
given the fact that the emission saturates in the CDS detectors.


The eastern parts of the ribbons (areas C, D, LP in the CDS image
and 1, 2 in the H$\alpha$ image) which are closest to the onset
location of the flare are not observed at the same time,
so a comparison of the velocities is problematic.
For the H$\alpha$ ribbons (region 3), we measure a strong downflow.
This could correspond to the impulsive downflow frequently observed
during the impulsive phase. But according to the timing, region
3 was observed at 15:34~UT (3 minutes after the start time of the flare). 
One possibility is that the flare was composed of more than one 
burst. Another possibility is the formation of a siphon flow inside a loop
joining the two areas  anchored in the two opposite polarities where the
main flare occurs.
This possibility is supported by the presence of opposite 
flows  (areas C and D) as seen  at 15:40-15:43 UT in the coronal lines.
(in Si XII and Fe XIX -55 km/s and  +20 km/s respectively). 
It is reasonable that  the evaporation was more effective in one
footpoint  and a siphon flow started a few minutes
after the flare.
Spectral observations with higher spatial resolution than 
CDS are needed, to study these kind of compact flares.

%We apparent that the ribbon plasma is very dynamic and it is
%difficult to have results completely consistent with the actual models.

%Perhaps future observations should be done  with a fixed slit waiting for the
%flare as did Brosius \& Phillips 2004 (ApJ, 613, 580), but it is difficult to get 
%H$\alpha$ spectra at the same place and the same time.




During a coordinated observing campaign, a compact and 
short-duration M1.0 flare was observed from the
ground and in space using spectrographs. 
In this paper and in the previous one (\citealt{berlicki_etal:04}),
which discussed the overall  morphology and the
magnetic topology of the active region, we adopted a multi-wavelength 
approach, i.e. we gathered together all the observational
information relevant for a detailed study of the flare.

Firm conclusions on the mechanisms that trigger the event
and detailed flare modelling cannot be achieved, given the 
(relatively) low spatial and temporal resolution of the current 
instrumentation. For example, as seen in the X-rays (GOES), 
the flare started at 15:31 UT and reached peak emission in just
3-4 minutes, i.e. on times shorter than the 
rastering times of the spectroscopic instruments.
The  SOHO/CDS and the
VTT/MSDP instruments do allow us to make 
detailed studies of e.g. 
the velocities in the ribbons over a wide range of 
temperatures (10$^4$ K to 8 $\cdot$ 10$^6$ K).
However, due to the scanning method of the spectrographs, 
we were only able to obtain spectra from a portion of the ribbons
at the peak of the X-ray emission, and from the central area 
of the flare during the decaying gradual phase.
Nontheless we are able to make some accurate measurements
of velocities over a range of temperatures, for which we
have proposed an explanation.

Other authors (e.g. \citealt{brosius_phillips:04}) 
have adopted a different observing approach, i.e. to fix the slit 
position and rely on the imaging instruments to study the 
morphology of an event. The slit was fortitously positioned
right where the flare took place.
%Obviously, neither of the two approaches is entirely 
%satisfactory. 

%Although only a few detailed studies 
%have been published, we can use them to provide some suggestions
%for future studies/observations.

Our observing campaign and analysis emphasises the importance of
combining chromospheric and coronal observations in order to
understand the dynamics of a solar flare. It is also essential to use
spectroscopic observations to obtain an unambiguous determination of
the plasma properties. Even so, with this set of observations, we have
only just glimsed the complexity of chromospheric evaporation. The main
problem with scanning instruments like VTT/MSDP and CDS is to be in the right
place at the right time! Neither the scanning approach or the
sit-and-stare approach is entirely satisfactory.
We look to future projects, such as SOLAR-B, which
will enable us to study active regions and flares in much more detail
through all layers of the solar atmosphere.

\begin{acknowledgements}

GDZ and HEM acknowledge support from  PPARC. 

This research was partially supported by the European Commission
through the RTN programme ESMN (European Solar Magnetism
Network, contract HPRN-CT-2002-00313).  AB and BS are
members of ESMN.

SOHO is a project of international collaboration between ESA and NASA.
 These observations were performed during a MEDOC
campaign (JOP 157).

We acknowledge the use of the SURF, TRACE, SOHO  
on-line databases.


\end{acknowledgements}



\bibliographystyle{aa}
%\bibliography{bib}
\bibliography{paper7}


\end{document}

