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1 Plasma flows during solar flares A. Berlicki Observatoire de Paris, Section de Meudon, LESIA (Astronomical Institute of the Wrocław University, Poland)

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Presentation on theme: "1 Plasma flows during solar flares A. Berlicki Observatoire de Paris, Section de Meudon, LESIA (Astronomical Institute of the Wrocław University, Poland)"— Presentation transcript:

1 1 Plasma flows during solar flares A. Berlicki Observatoire de Paris, Section de Meudon, LESIA (Astronomical Institute of the Wrocław University, Poland)

2 2 Analysis of the vertical plasma flows manifested in shifts or asymmetries of spectral lines Evaporation processes during solar flares This work: for better understanding these phenomena: a) spatial distribution of evaporating regions within the whole flare ribbons is not clear, b) the location of these regions with respect to X-ray loop footpoints where the heating can occur is also not well evidenced, c) the velocities of the evaporative plasma are not yet well determined; d) no information about the vertical structure of the velocity field in evaporative areas.

3 3 1) EXPLOSIVE CHROMOSPHERIC EVAPORATION - impulsive phase TWO TYPES OF EVAPORATION OBSERVATIONAL EVIDENCES: a) large blueshift observed in hot coronal and TR lines (Ca XIX, Fe XXV)  upflow of hot (T ~ K) plasma with V > 100 km s -1 (Wulser et al., 1994) - Yohkoh/BCS, SOHO/CDS b) redshift of chromospheric spectral lines - downflow of cool plasma with V > 20 km s -1 (Fisher et al. 1985; Fisher 1987; Švestka 1976; Ichimoto & Kurokawa 1984) - ground base observations Energy deposition layer UPFLOW - hot lines DOWNFLOW - cool lines

4 4 THEORY: a) driven mainly by non-thermal electrons accelerated during the primary energy release (Antonucci et al. 1984), at the initial stage of solar flares; thermal conduction front - alternative driving mechanism b) occurs 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 (Fisher 1987). c) the temperature of the heated region is above 2x10 5 K and thus this kind of evaporation cannot be seen in the chromospheric spectral lines. d) at the same time the low temperature chromospheric plasma is compressed and ‘pushed’ down - ‘chromospheric condensations’ TWO TYPES OF EVAPORATION 1) EXPLOSIVE CHROMOSPHERIC EVAPORATION - impulsive phase

5 5 TWO TYPES OF EVAPORATION 2) GENTLE CHROMOSPHERIC EVAPORATION - gradual phase OBSERVATIONAL EVIDENCES: a) small blueshift observed in coronal and TR lines (SOHO/CDS)  upflow of hot plasma with V < 80 km s -1 (Czaykowska 1999; Brosius 2003) b) chromospheric spectral lines exhibit also blueshift - upflow of cool plasma with V < 15 km s -1 ( Schmieder et al. 1987). Schmieder et al. 1987

6 6 THEORY: a) driven by the large conductive heat flux from a high-temperature flare plasma contained in magnetic tubes above the chromosphere. (Antiochos and Sturrock 1978; Schmieder et al. 1990; Czaykowska et al. 1999) b) appears during the gradual phase of solar flares, when there is no significant flux of non-thermal electrons. c) the temperature of the heated region is around 10 4 K and this kind of evaporation can be observed in the chromospheric spectral lines. d) other mechanism - decompression of chromospheric plasma after the impulsive phase (???) 2) GENTLE CHROMOSPHERIC EVAPORATION - gradual phase

7 7 To understand the evaporation mechanisms we use the multiwavelengths spectroscopic observations of solar flares in: 1) Optical range: analysis of the chromospherical line profiles and using the non-LTE radiative transfer codes  we obtain the models of flaring atmosphere and the velocity field during solar flares 2) EUV range: analysis of the EUV TR and coronal lines (SOHO/CDS)  physical conditions and velocity field in higher layers of the flaring atmosphere, energy transport 3) X-ray range: spectroskopic observations (RHESSI) in the energy range E > 3keV  analysis of thermal and non-thermal processes during the different phases of solar flares, mechanisms of chromospheric heating driving the evaporation

8 8 My work... An application of the non-LTE radiative transfer methods in the modelling of the gentle evaporation processes during the gradual phase of October 22, 2002 flare Data: spectroscopic observations of H  line (VTT/MSDP - Tenerife)

9 9 N THEMIS/MSDP 15:41-16:29 UT, B LONG VTT/MSDP, H  AR /1. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - overview of the flare Data: 22-Oct-2002, M1.0 flare at 15:29 UT Berlicki et al., A&A, 423, 1119 (2004)

10 10 The lfff extrapolation of the magnetic field explain the temporal evolution of the H  flare ribbons: R1 - compact and with constant position, R2 - elongated and exhibited fast propagation along itself, from the emerging bipole towards the leading sunspot, Propagation of R2  explained by the reconnection of the growing emerging field lines with a higher and higher overlying field. Huairou Vector Magnetogram THEMIS B LONG R2 SHEARED EMERGING BIPOLE lfff 1/2. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - overview of the flare

11 11 QS Area VTT/MSDP 16:00 UT 1/3. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - gradual phase - observed profiles Asymmetric profiles at 16:00 UT Photometric calibration: using the mean QS profile and the reference profile (David) Average signal from nine pixels inside the square box 0.75’’ x 0.75’’ To process the MSDP data the code of P. Mein was used To analyse the H  line profiles we chose six areas:

12 12 1/4. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - gradual phase - observed profiles These six areas were analysed at six times (gradual phase)  36 profiles from VTT/MSDP within the spectral range 6563  0.9Å Time evolution of the integral intensity in the H  line within  0.7 Å from the line centre. In all areas we observe slow decay of the intensity  gradual phase

13 13 1/5. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - method We used observed H  line profiles to construct the models of solar chromosphere Observed profiles were compared with the grid of synthetic H  line profiles calculated by the non-LTE code developed by Heinzel (1995): Assumptions: * 1D plane-parallel geometry (H CHROM << L CHROM & H CHROM << R SUN ) * hydrostatic equilibrium (long time-scales during the gradual phase) * MAVN - F1 semiempirical flare model of Machado et al. (1980) as the reference model. Other models (VALs, FALs etc.) are also possible. Non-LTE code: allows us to find a model of the flare atmosphere * two imput parameters: m 0 and  T, * macroscopic velocity introduced to analyse the asymmetries or shifts of the line profiles 

14 14 1/6. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - method To investigate the influence of the macroscopic velocity field we adopted the formula used by Mihalas (1978) to describe the expanding atmospheres:  m - optical depth in the atmosphere where the velocity has a value V O we used the optical depth in the H  line centre For  <<  m (upper layers of the chromosphere): V  2V0 For  >>  m (close to the photosphere) : V  0

15 15 1/7. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - simulations The distribution of the velocity in a flare atmosphere for different values of  m (left panel) and for different values of V 0 (right panel).

16 16 V O = -8 km s -1 (left) and +8 km s -1 (right) log  m = 0.1: ··········· log  m = 1. 0 : log  m = : ——— Simulation of influence of the velocity field on the emitted H  line profile: 1/8. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - simulations QS The value of the velocity at the depth of H  formation has a significant influence on the shape of the profile

17 17 H  contribution function ( CF ): 2D function of the wavelength from the line centre and the depth in the atmosphere; CF - represents the depth of the spectral line formation 1/9. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - contribution function Kašparová & Heinzel, 2002

18 18 1/10. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - construction of the grid HOW DO WE FIND THE MODELS ? FIRST STEP: grid of static models - by modifying two input parameters m 0 and  T: m rm - column mass at all layers ID of the reference F1 model T rm - temperature of the reference F1 model m 1 - column mass where we start to modify the temperature Each observed profile was fitted by the least-square technique with all synthetic profiles from the „static” grid  we determined the static models for all analysed areas To construct the grid we used 23 different values of m 0 and 26 values of  T (598 models)

19 19 1/11. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - construction of the grid SECOND STEP: grid of models with velocity - by introducing the velocity field defined as: Finally, for each H  profile observed in the analysed areas of the flare, we found the best fit among the synthetic profiles  we found the semiempirical models and the height distribution of the velocity within all analysed areas.  m - optical depth in the atmosphere where the velocity has a value V O We used 21 different values of  m and 33 values of V 0 (693 models) Again, each observed profile was fitted with all synthetic profiles from this grid with velocity  we found the models with the velocity for all analysed areas Why two-steps fitting ? Since the velocity is not large we could do this instead of fitting models !

20 20 1/12. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - results Velocity distribution in the area ‘5’ ········ - reference QS profile  - observed profile – – – - fitted synthetic profile

21 21 1/13. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - conclusions RESULTS AND CONCLUSIONS 1) For the first time the evaporative flows in the gradual phase are studied quantitatively by using a non-LTE radiative transfer code and the models with the velocity field. Most of the profiles obtained in the analysed areas of the flare ribbons exhibit a red asymmetry. As we could see from the theoretical calculations in the case of self-reversed profiles the red asymmetry means an upflow.

22 22 1/14. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - conclusions 2) In most ot the analysed areas the time dependence of the velocity does not show any significant changes of sign.

23 23 1/15. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - conclusions 3) The quality of fitting the synthetic profiles with the observed ones is different for the different areas and times. High values of  2 are associated with the profiles, the shape of which cannot be reliably reproduced by our codes. Nevertheless, most of the profiles fitted with reasonable accuracy (low  2 ) exhibit evident predominance of upflows.

24 24 1/16. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - conclusions 4) We interpret the upflows found in the flare ribbons in terms of the Antiochos and Sturrock (1978) model for gentle evaporation. This process can be driven by conductive heat flux.

25 25 1/17. Observed H  line profiles and the non-LTE codes - velocity field and flaring chromosphere models - future IN THE FUTURE It would be interesting to use more spatial points at more times and to use the spectra obtained within a wider range of wavelengths from the line centre. Other distributions of the velocity field in the chromosphere should also be tested. In addition, to perform non-LTE modelling of the flare structure it would be useful to have other spectral lines formed at different levels of the chromosphere like hydrogen H , H , infrared calcium line (Ca II 8542Å) etc.

26 26 THE END


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