1. Changes in the spectral line profiles during a solar flare
Abramenko, V.I.
Big Bear Solar Observatory, NJIT
And
Baranovsky, E.A.
Crimean Astrophysical Observatory, Ukraine

2. Modeling of spectral line profiles:
INTRODUCTION
CaI 6103A?
Magnetographs which measure the magnetic field in solar flares suffer from ambiguities caused by changes in the spectral line profile during the flaring process.
This problem creates difficulties in the interpretation of flare-related changes in the photospheric magnetic field.
Observations of spectral line profiles in flares:
The study presented here was encouraged by a very frequent question of referees on any our results concerning to changes in the magnetic field during a solar flare. Usually, a referee says: “The observed changes in the magnetic field during a flare may be an artifact caused by changes in the profile of the spectral line used in your magnetograph.” So, we decided to clear up this question,,, to study, what kind of changes in the profile of the spectral line one should expect when observing the magnetic field during a flare.
First of all, we found out that publications on the spectral line profile measurements during solar flares are very scanty.The same may be said about the modeling of flare-related spectral line profiles. Only the profile of the Nickl I 6768A line was modeled under the flaring conditions because this line is used in SONO/MDI. The result of observations and modeling can be formulated as follows: during a flare, an enhancement of line core intensities is observed and it can be explained by heating of the photosphere.
When a flare is very strong, then the emission can arise in the line core, which may result in an appearance of so called magnetic transient: small patch of inverse polarity. ((This effect was observed at first time by Zirin and Tanaka for the FeI5323 spectral line. ))
But what about our CaI 6103 spectral line? Let us have a look at two magnetograms taken before a flare and at the maximum of the X1.2 flare on October 22, 2001.
Modeling of spectral line profiles:  I

3. BBSO/DMG data, October 22,2001, the X1.2 flare
Before the flare
During the flare
This is the magnetogram obtained before the flare, and the second one - at the flare maximum. The third panel is the H-alpha image of the flare. This(#1) is the strongest kernel of the flare. Before the flare there was a negative polarity (#2), (#3) there, but at the flare maximum we see (#4),(#5) an area of positive polarity on the place of the brightest H-alpha kernel. This is a so called “Magnetic Transient”. ((So, we may suspect that our spectral line, Calcium 6103, also may be affected by a flare influence, namely, in the line core absorbtion may be replaced by emission in a place of very strong brightness.)) Let us see what happened in a place of a moderate brightness. (#6).This is (#7),(#8, thick line) the magnetic field before the flare, and – at the flare maximum (#9-ov),(#10-dotted line). We see that the peak magnitudes of the magnetic field (of both polarities) are reduces(#11-%), (#12 -%), (#13-%),(#14-%) by 12-40%.
Are such changes real or they are artifacts caused by the flare influence? To answer this question, we have to analyze the profile of the Calcium 6103 spectral line measured during a flare. We failed to find observations of this kind in the literature.
12%
40%

4. Observational Data: 2N flare on Oct 11,1960 Slit position
We were lucky to find in archives of the Crimean Astrophysical observatory the spectral observations of a moderate flare. One may say that this is a Pre-historical flare: there were no GOES, no SGD at that years, we can say only that it was a moderate flare of optical class about 2N. The spectra were obtained with the Crimean Eshelle spectrograph – an instrument which allows us to simultaneously obtain the most part of the Fraungofer spectrum of the Sun on a photo-plate. The slit position is shown here – across one of the flare ribbons.

5. Measured spectral line profiles:
Solid lines – in the flare,
Dotted lines – in a quiet-
sun area
Using a modern densitometer, we measured spectral line profiles for six lines. The flare line profiles are shown here by solid lines.
We selected then another photo-plate there a quiet-sun area was recorded at the same distance from the disk center as the flare was. And we measured the quiet-sun profiles . They are shown here by dotted lines. One can see that the flare profiles noticeably deviate from the quiet-sun profiles for all of the spectral lines. A residual intencity at the core of 5 lines appears. In addition, for the 4 spectral lines, the slope of the wings is less steep in the flare profiles.
We payed special attention to the Ca 6103 spectral line (#1) because this is the line which is used by our magnetograph. Let us consider this panel in details.
BBSO/DMG: CaI 6103A
Spirock et al. 2001

6. dI/d Quiet-sun profile Delbouille Atlas profile
(#1) – this is a quiet-sun profile measured from the crimean suiet-sun photo-plat.
(#2) – this is a quei-sun profile of the Ca 6103 spectral line from [Delbui] Atlas.
(#3) – The vertical dashed lines show the BBSO calibration range(according to e paper by Varsik in Solar Physics, 1995).
This is the linear range where the slope of the line profile (#4), (#5) [di-ai-bai-di lambda], was determined for the slope calibration method ….(# next slide…)
BBSO/VMG
Calibration
range

7. The parameter of calibration determined from quiet-sun data
Varsik, Solar Phys. 161 (1995):
According to the profile slope method for the calibration
of the longitudinal magnetic field, the parameter V
is obtained by subtracting one pair of frames taken at
alternating polarities of the electro-optic modulator.
V is then related to the magnetic field strength, B, as
To be determined
Constant
… discribed in details by John Varsik. ((Dal’she vse chitai po slaidu)).
The calibration routine is reliable when there are no flaring in an active region.
Measured
The parameter of calibration
determined from quiet-sun data

8. In the flare the slope dI/d decreased by about 20%
profile
But what happens during a flare?
(#1), (#2) – this is the profile of our spectral line during the flare (measured from the crimean photo-plate). The most interesting question for us is what happens in the in the blue wing,, in the calibration range?
(#3), (#4) – we see that the slope, di/d lambda, decreased in the flare…
(#5) – bu about 20% for this flare…
(#6 – str) – which means that (chitai: using the existing routine for calibration (with the slope determined for a quiet sun) , we underestimate the magnitude of the longitudinal field by about 20% for a moderate flare.)
Using the existing routine
for calibration (with dI/d
determined for a quiet Sun),
we underestimate the magnitude
of Bz by about 20% for a
moderate flare.

9. A simiempirical modeling
A model of the photosphere
for the 2N flare on Oct 11, 1960
A simiempirical modeling
Of the photosphere by
Baranovsky, E,A. and
Shoumko, A.V. 1999
To propose a possible explanation of the observed differences between flaring and quiet-sun profiles, we performed a semiempirical modeling of the flaring photosphere. Namely, we specified functions of temperature, gas pressure, density, magnetic field , - with height in the photosphere, substituted these functions into the radiation transfer equation, and, solving this equation, we obtained spectral line profiles for our six spectral lines. Then we compared the calculated profiles with the observed ones. When the agreement between the calculated and the observed profiles is not good, the calculations repeated for another set of model functions, and so on until the proper agreement of the profiles is obtained. So, it is kind of iteration method. It is widely known tecknique , and the calculated spectra are referred to as synthetic spectra.
The top panel shows the functions for the temperature and magnetic field which gave us the best agreement between the observed and the calculated line profiles for all six lines.

10. The flare profiles:
Observed – solid lines;
Modeled – dotted lines

11. A model of the photosphere for the 2N flare on Oct 11, 1960
Our experience in the modeling told us that a hot layer in the upper photosphere is responsible (#1) for the enhanced residual intensity (#2) in the line core, whereas inhomogenity of the magnetic field in deeper layers may produce the decrease(#3) of the wing slope during a flare. This inference is also supported by analysis of the contribution functions : This function shows which layers in the photosphere contribute into the emission(or absorbtion) at a given spectral range. We see that the core of the line is formed higher and closer to the hot layer than the wings. The wings form deeper and farther from the hot layer. the emission in wings may be affected by different factor, such as inhomogeneity of the magnetic field in the photosphere among others.
This result allows us to assume that there exists a flare-related enhancement of the complexity of the photospheric magnetic field, which is in qualitative agreement with previous study of flare-related changes in the photospheric magnetic field.
It is not excluded, however, that other mechanisms, like the increase of the turbulent veloscity, might also contribute into the line profile broadening, as it was observed in the chromosphere and in the corona.

12. Conclusions Measuring the longitudinal magnetic field in the areas
of bright H ribbons at the flare maximum using the
CaI 6103A spectral line and the profile slope calibra –
tion method, we underestimate the magnetic flux of both
polarities by about 20% for a moderate flare.
The amount of reduction depends on the magnitude and phase of a flare, on the position in flare ribbons, and so on. This question deserves further analysis on the basis of new and more extensive observational data.
We are thankful to John, Tom, Vasyl, Alla and Alex for very useful discussions and suggestions .