WDS'06 Proceedings of Contributed Papers, Part III, 27–32, 2006.
ISBN 80-86732-86-X © MATFYZPRESS
Asymmetries of Spectral Lines in Solar Flares
T. Prosecký
Charles University, Faculty of Mathematics and Physics, Astronomical Institute, Prague,
Czech Republic.
Academy of Sciences of the Czech Republic, Astronomical Institute, Observatory
Ondřejov, Czech Republic.
Abstract. This contribution deals with line profile asymmetries in solar flares.
Basic knowledge of asymmetries and review of possible explanations and recent
works is presented here.
Introduction
Solar flares belong to the most interesting solar activity phenomena. It is a process during
which a lot of energy, deposited in upper parts of the solar atmosphere, is released suddenly.
Detailed explanation of all flaring processes still does not exist; one of current problems is the
analysis of vertical mass motions of a flaring plasma and determination of the velocity fields
and exact physical conditions which lead to these processes. One way, which can contribute to
confirmation or disapproving of the assumed scenarios is the analysis of line profile asymmetries
in solar flares.
Electromagnetic spectrum produced by solar flares has a very wide range of wavelengths,
usually from hard X-rays (wavelengths ≈ 10−2 nm) or sometimes even from γ-rays (wavelengths
≈ 10−4 nm, photons with energies exceeding 10 MeV) to radio waves (wavelengths up to ≈
104 m). The character of the radiation of a typical solar flare is a thermal one, only at the very
short wavelengths (hard X-rays with λ ≤0.1 nm) and in the range of very long radio waves a
short time non-thermal outbursts (bremsstrahlung, synchrotron processes) can be seen. This
is caused mainly by beams of the non-thermal particles with high energies [Tandberg-Hanssen
& Emslie, 1988]. Different parts of spectra originate at different layers of the solar atmosphere
and depend (but not only) on the temperature of the corresponding layers - we speak about
Hα , X-ray or EUV flares. It can be seen that the analysis in flares on several wavelengths can
reveal the physical conditions in various parts of the solar atmosphere.
Optical spectrum of solar flares is characteristic by the emission lines, especially lower
lines of the Balmer series and the lines H and K of the ionized calcium. Profiles of these lines
are usually very broad (up to 1 nm or even more) and various Doppler shifts and the so-called
asymmetries can be observed as well, the situations when one line wing is wider and more intense
than the other one. We can speak either about a red asymmetry (radiation in the red wing
dominates) or a blue asymmetry (the blue wing is dominant). Such line profile asymmetries
were recorded for the first time by Waldmeier [1941].
Analogous shifts and asymmetries can be found at other wavelengths, important are mainly
EUV and soft X-ray parts of spectra (see below). Spectral lines visible in the UV or EUV range
are predominantly formed at temperatures between 104 K up to ∼ 107 K, so they can be used
for an analysis of all parts from the upper chromosphere to the solar corona [Varady, 2002].
‘Standard model’ of solar flares and possible explanation of line profile
asymmetries
Now it seems to be clear that the source for the energy released in solar flares is in magnetic
fields. These fields are generally stable in the upper chromosphere and inner corona. However
a deformation of pre-existing magnetic field lines can take place here as a consequence of the
convective mass motions under the photosphere. This leads to a growth of the gradient of
magnetic induction vector and therefore to induction of currents and creation of current sheets.
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PROSECKÝ: ASYMMETRIES OF SPECTRAL LINES IN SOLAR FLARES
Figure 1. Lower left panel: Example of the record of flare in NOAA 9393, April 2, 2001.
Taken by Ondřejov MFS device that recorded observations in Hα, Hβ and infra-red Ca II 8542
Å lines simultaneously. So-called slit-jaw image (record of a flare in Hα) can also be seen. See
the text for details. Upper right panel: Example of Hα line profile with strong red asymmetry
(impulsive phase of the flare, obtained by the reduction of the lower left picture). In the picture
also the quiet Sun profile (dashed line) and tabulated quite Sun profile (dotted line, [David,
1961]) can be seen, both used for reduction of observation from instrumental to real physical
quantities. Upper left and lower right panel: Profiles of Hβ and Ca 8542 lines respectively, also
obtained by the reduction of the lower left image. All lines exhibit strong red asymmetries.
Final effect is a transformation of coronal magnetic field into higher energy states.
At a certain moment, usually due to some instability related to the current sheet, the
coronal magnetic field starts changing its configuration to the most preferable state via a process
called reconnection of the magnetic fields. This is usually called preflare phase, it takes several
minutes and can be detected in extreme UV and soft X-ray range of spectra [Varady, 2002].
The so-called impulsive phase is started by the reconnection. It comes due to deposited
energy in a form of the electromagnetic radiation, heating of plasma (at the place of reconnection) and mainly in the form of a kinetic energy of the so-called particle beams. The beams
of accelerated (non-thermal) particles propagate (along new configuration of the magnetic field
lines) to lower parts of the solar atmosphere, interact with particles of the background plasma
and dissipate their energy at these layers. The background plasma is heated by this energy and
a brightening in strong chromospheric lines like e.g. Hα can be clearly observed. It leads to
fast changes of intensities in many lines during the impulsive phase.
It is also widely accepted that the heated plasma is evaporated into the coronal part of the
solar atmosphere during chromospheric flares providing the material for the so-called post-flare
loops, system of hot loops visible in soft X-rays for hours [Antiochos & Sturrock, 1978 ]. This
scenario is supported mainly by soft X-ray observations - such evaporation and mass motions
should produce blue-shifted line profiles that are really observed at temperatures of several
MK (temperatures of coronal structures can reach up to 30 MK [Varady, 2002]) by the Yohkoh
observations [Berlicki et al., 2002] or by the Coronal Diagnostic Spectrometer (CDS) on board
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PROSECKÝ: ASYMMETRIES OF SPECTRAL LINES IN SOLAR FLARES
Figure 2. Right panel: Example of Hα line profile with a strong blue asymmetry (impulsive
phase of the flare, obtained by the reduction of the image in left panel). On this picture one
can see also the observed quiet Sun profile (dashed line) and tabulated quiet Sun profile (dotted
line, [David, 1961]). The asymmetry had been observed at the same active region as the one in
Figure 1 just only one hour later. Compare with Figure 1.
of SOHO [Teriaca et al., 2003]. At the end of the impulsive phase the plasma starts to cool
down (radiatively), become denser and falls down back to the chromosphere [Varady, 2002].
Subsequent decrease of the intensity can take several hours that is called the gradual phase of
a flare. This is also well observed in the Hα line.
In the last years the existing solar flare models predict two kinds of the chromospheric
evaporation processes. First type is referred to as gentle evaporation [Antiochos & Sturrock,
1978 ] and this could be observed in chromospheric lines (e.g. Hα, Ca II 8542 Å) or in EUV lines.
This type of evaporation takes place when the chromosphere is heated either by non-thermal
particles with low flux (≤ 1010 erg cm−2 s−1 , Fisher et al., [1985]) or by thermal conduction
in later phases of a flare (however the mechanism is still debated, [Berlicki et al., 2004]). Final
effect is a slow upward hydrodynamic expansion with velocities of about ten km/s. Second
type is referred to as explosive evaporation. This takes place when the flux of non-thermal
particles is high enough (≥ 3.1010 erg cm−2 s−1 ) and the chromospheric plasma is not able to
radiate energy at a sufficiently short time (the heating rate by the beams is bigger than the
radiative losses). This leads to an expansion with velocities of about hundreds of km/s into the
overlaying layers. An overpressure of the evaporated material also drives a low-velocity (tens
of km/s) downward mass motions of the cool and dense material into the lower layers of the
chromosphere. This process is the so-called chromospheric condensation [Antonucci et al., 1984,
Fisher et al., 1985].
Now it seems to be clear that the process of the chromospheric condensation is responsible
for red shifts and asymmetries observed by many authors [e.g. Ichimoto & Kurokawa, 1984;
Wülser & Marti, 1989; Canfield et al., 1990] in chromospheric lines. For example of such
asymmetry see Figure 1. The whole scenario is also supported e.g. by Brosius & Phillips [2004]
who have recorded up-flows and down-flows simultaneously during impulsive phase of a flare.
However, problems arise with the explanation of the blue asymmetry (for example see
Figure 2) that is sometimes observed at the beginning of the impulsive phase of some flares
during several minutes, see e.g. Canfield et al. [1990] or Heinzel et al. [1994] where also a
review of older observations of blue asymmetries can be found. First explanations of the origin
of the blue asymmetry were in terms of an upward mass motion (scenarios of localized heating of
the deep chromosphere, heating of upward-moving cool material or magnetic-field line squeezing,
see Canfield et al. [1990] for details), opposite of the origin of a red asymmetry. However, as
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PROSECKÝ: ASYMMETRIES OF SPECTRAL LINES IN SOLAR FLARES
was shown in Heinzel et al. [1994], the real situation can be different. Blue asymmetry can
origin due to downward mass motions - because of the absorption in the red wing of a spectral
line in a cool and dense material. As a consequence the blue wing of the line has relatively
higher intensity than the red one. For details see the scenario of an electron-beam heating with
the return current in Heinzel et al. [1994] .
In summary, it can be seen that the situation can be rather complicated and still is not
clear. For this reason new observations and a more precise analysis should be carried out.
Aim of investigation
Aim of my study should be a detailed analysis of line profile asymmetries in solar flares, using a rich archive of the Ondřejov Multichannel Flare Spectrograph (MFS) device that was in use
till June 2004, and new observations obtained by reconstructed Horizontal-Sonnen-ForschungsAnlage (Horizontal device for solar research, HSFA2).
MFS
This device was in use for almost five decades. Many informations about the original
construction of the device can be found in Valnı́ček et al. [1959]. MFS was constructed for
observations of solar flares or prominences simultaneously in selected - diagnostically important
- spectral lines. At the beginning of 90’s the device has been reconstructed [Kotrč et al., 1993].
The CCD detectors with dimensions 6.4 mm and 4.8 mm and working frequency 25 exposures
per second were installed. A typical data output can be seen in Figure 1. MFS device recorded
observations in Hα, Hβ and infra-red Ca II 8542 Å lines together with the so-called slit-jaw
image - the record of a selected flaring part of the Sun (in Hα line) across which the slit of
the spectrograph is placed. Because of this it can be seen clearly which part of the Sun is
responsible for the recorded spectral profiles. Final image was recorded at video cassettes.
HSFA 2
It is possible to find an information about this modern device at the web site of the Astronomical Institute, Academy of Sciences of the Czech Republic, see references. The spectrograph
can record the spectrum of Hα, Hβ, He D3, Ca II H or K line and the slit-jaw image using 5
CCD cameras. There are several advantages against old MFS device: better resolution, high
quality optics, better seeing and a direct record in digital form. This modern set-up gives us
a chance for a more detailed and precise analysis of our future observations. Example of spectrum is in Figure 3, compare with Figure 1. Detailed information can be also found in Kotrč &
Kschioneck [2003].
Figure 3. Example of an (Hα) spectrum taken by the HSFA2 device.
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PROSECKÝ: ASYMMETRIES OF SPECTRAL LINES IN SOLAR FLARES
Methods and techniques of analysis of line profile asymmetries
As was shown above there are many possibilities of the plasma flows that can cause different types of asymmetries. There are two basic methods that can be use for the analysis of the
line profile asymmetries and the calculation of vertical velocity fields. First one, the bisector
method, has been widely used in earlier studies, but there are several important disadvantages
of the method like its straightforwardness - the origin of asymmetries could not be found unambiguously - or an allocation of the derived velocities to geometrical or optical depths in the
flaring atmosphere. More details can be found in Prosecký et al., [2006].
Second method is the modeling of the radiation transfer in a flaring atmosphere and comparison of the observed and synthetic non-LTE profiles. Basic information about the radiation
transfer and non-LTE models of stellar atmospheres can be found for example in Hubeny [1996].
To study the line profile asymmetries it is necessary to include the velocities into the basic nonLTE calculations. In the case of small velocities (≤ 10 km/s) which do not affect the plasma
excitation and ionization too much it is possible to use a static non-LTE models for finding
the atmospheric parameters and then to perform the formal solution of the radiation transfer
equation with a prescribed velocity field [Nejezchleba, 1998]. Here our analysis of line profile
asymmetries will start. We will use this approximation and will follow the approach of Berlicki
et al., [2005] at first. Next steps will include a modification of the basic non-LTE calculation
code (see [Heinzel, 1995]) for general velocities. Basic statistical data analysis (percentage of
the occurrence of asymmetries, correlation of optical, EUV and X-ray intensities) will be also
performed.
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