Hinode/EIS Nugget – Velocity Characteristics of Evaporated Plasma

by Ryan Milligan, GSFC



Chromospheric evaporation is largely accepted to be the process by which solar flares achieve their high temperatures and densities. The standard flare model states that electrons are accelerated at or near a magnetic reconnection site in the corona and then travel along newly reconnected magnetic field lines toward the chromosphere. Here, they are decelerated by the increasingly dense atmosphere and typically lose their energy by one of two mechanisms: an encounter with a proton or ion will result in the emission of a HXR photon through the bremsstrahlung process; Coulomb collisions with ambient electrons, on the other hand, result in an overall heating and, by consequence, expansion of the chromospheric material. The velocity at which the evaporated material rises is dependent on the total energy flux of accelerated electrons (in erg cm-2 s-1) that reach the footpoint and has traditionally been measured through Doppler shifts of EUV and SXR emission lines. EIS provides the capability to determine the location and magnitude of the evaporation flows at high spatial, spectral, and temporal resolution for a myriad of emission lines formed over a wide range of temperatures simultaniously, marking a significant advancement over previous instruments. Coordinated HXR observations from RHESSI can establish the location and the parameters of the driving electron beam needed to test the predictions of flare models.


By January 2008 Hinode had lost the use of its X-band transmitter. As such, the amount of data being telemetered to the ground was severely reduced. Also, in November 2007, the detectors on RHESSI had been successfully annealed after five years of radiation damage, bringing them back to specifications comparable to early 2005. Therefore, December 2007 marks a unique time frame when both Hinode and RHESSI data sets were near optimal. Fortunately this was also a fairly active period of solar activity with several flares observed by both instruments. The event presented here was a C1.1 flare that occurred on 14 December 2007. The left hand panel of Figure 1 shows the active region in which it occurred as observed with TRACE. The box inset shows a zoomed in portion where the flare occurred with the two loop footpoints clearly visible. The corresponding HXR emission from RHESSI also aligns with this footpoint emission. The right hand panel shows the RHESSI lightcurves in 3 energy bands.

Figure 1: Left panel: NOAA 10978 as observed by TRACE during the impulsive phase of the flare. The inset shows the footpoint emission during the impulsive phase with the 20-25 keV emission observed with RHESSI overlaid. Right panel: RHESSI lightcurves in the 3-6, 6-12, and 12-25 keV energy bands. The vertical solid line marks the time of the TRACE image. » Click figure to see full-size image.

Figure 2: Left panel: An Fe XVI raster taken during the flare's impulsive phase. Right panel: The corresponding velocity map. The contours overlaid denote the 20-25 keV emission as observed with RHESSI. » Click figure to see full-size image.
The observing study that EIS was running when the flare occurred (CAM_ARTB_RHESSI_b_2) was originally designed to search for active region and transition region brightenings in conjunction with RHESSI. Using the 2'' slit, EIS rastered across a region of the Sun from west to east covering an area of 40''x144'', denoted by the rectangular box in Figure 1. Each slit position had an exposure time of 10 s resulting in an effective raster cadence of ~3.5 minutes. These fast-raster studies are preferred for studying temporal variations of flare parameters while preserving the spatial information. Equally important though is the large number of emission lines that covered a broad temperature range. This observing study used 21 spectral windows, some of which contain several individual lines. The work presented here focuses on 15 lines spanning the 0.05-16 MK temperature range. The majority of these lines are well resolved and do not contain blends, thereby reducing ambiguities in the interpretation of their analysis. An example of one of the rasters taken during the impulsive phase is shown in Figure 2 (Fe XVI, 2.5 MK) along with its associated Doppler velocity map. Again, the footpoints are visible near the top of the raster and align well with the HXR footpoints as observed by RHESSI. The velocity map shows that the bright footpoint emission at this temperature is blueshifted (denoting upflowing plasma) as expected.

Not Expected

Many previous studies that investigated the chromospheric evaporation process relied on detecting blueshifts in a single (or small number of) emission line(s). In this event blueshifts were detected in 7 different ionization stages of iron (Fe XIII, Fe XIV, XV, XIV, XVII, XXIII, and XXIV; i.e. >1.5 MK. See Figure 3). A linear relationship was also found between the magnitude of the velocity and the formation temperature of the line, confirming theoretical predictions which state that hotter material should rise faster due to the pressure gradient relative to the overlying corona. However, it was also found that emission lines formed below 1.5 MK were redshifted, which was not expected. Early models (and indeed, observations) state that only material at chromospheric and transition region temperatures (<0.5 MK; i.e. below the site where the electrons are believed to deposit their energy) should be redshifted due to the overpressure of the rising material. These observations would therefore suggest that either the electrons become thermalized at higher (coronal) altitudes, or the chromosphere is being somehow ''backwarmed'' by the incident electron beam as it recoils. There is also a possibilty that these ''cooler'' lines were formed out of ionization equilibrium which could affect the assumed formation temperature.

Figure 3: Evaporation velocity as a function of temperature for each of the 15 lines used in this study. Negative velocities are blueshifted (upflowing) while positive values are redshifted (downflowing). The dashed lines are least-squares fits to the blueshifted and redshifted (excluding the He II line) data points. (Note that this is a log-normal plot so both least-sqaures fits are linear.) The values in the top left corner correspond to the parameters of the electron beam as measured by RHESSI at the time of the flows (delta: spectral index, Ec: low-energy cutoff, F: electron energy flux). » Click figure to see full-size image.

Perhaps an even more unexpected result of this analysis was that the line profiles taken from the footpoints for the two hottest lines (Fe XXIII, 12 MK and Fe XXIV, 16 MK) were both dominated by stationary components despite exhibiting upflows >250 km/s (Figure 4). Many flare heating models predict that any chromospheric material heated during the impulsive phase should rise due to the pressure gradient relative to the overlying corona. The corresponding line profiles should therefore be completely shifted, which was the case for all other lines in this event. Comparisons made with other, similar events will reveal whether this was a unique, one-off case (perhaps some unresolved magnetic structure confined the hottest material at rest) or whether there is a flaw in our understanding of the evaporation process.

Figure 4: Fe XXIII and Fe XXIV line profiles from the brightest footpoint pixel, within the HXR contours. The solid vertical line denotes the approximate rest wavelength of each line. Note that both line profiles are dominated by a near-stationary component while the dashed lines mark the fits to the highly blueshifted components. » Click figure to see full-size image.


Despite the recent quiescent period of solar activity EIS is providing a unique insight into the effects of electron beam heating in the chromosphere during solar flares. Coordinated observations between EIS and RHESSI during Cycle 24 can be used to diagnose the temporal variations of this beam heating, as well as nonthermal line broadening and density variations.

1. Milligan & Dennis (2009), ApJ, 699, 968

Contact the author: Ryan Milligan   Visit author's webpage

Last Revised: 1-Jun-2010

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