DIFFERENT PATTERNS OF CHROMOSPHERIC EVAPORATION IN A FLARING REGION OBSERVED WITH HINODE/EIS

Chromospheric evaporation refers to the drastic mass motions in flaring loops caused by rapid energy deposit in chromospheric layers (or probably higher) by non-thermal electrons or thermal conduction. When the flare energy is transported toward lower layers, it can produce a local overpressure that drives both upward and downward mass motions. These motions can be detected through Doppler shift measurements in chromospheric and coronal lines. Antonucci et al. (1982, 1985), Antonucci & Dennis (1983), Canfield et al. (1987), Zarro & Lemen (1988), Wülser et al. (1994), Ding et al. (1996), and Doschek & Warren (2005) obtained blueshifts of 200–400 km s⁻¹ in the Ca xix line using data from the Bent and Bragg Crystal Spectrometer (BCS) on board the Solar Maximum Mission (SMM; Acton et al. 1980) and Yohkoh/BCS (Culhane et al. 1991) data, respectively. Similar measurements using data from the Coronal Diagnostic Spectrometer (CDS; Harrison et al. 1995) on board the Solar and Heliospheric Observatory (SOHO) revealed upflow velocities of 60–300 km s⁻¹ in the Fe xix line (Teriaca et al. 2003, 2006; Brosius & Phillips 2004; Del Zanna et al. 2006; Milligan et al. 2006a, 2006b). On the other hand, in chromospheric and transition region lines, redshifts of around tens of km s⁻¹ were observed, implying downward motions (Wülser et al. 1994; Czaykowska et al. 1999; Teriaca et al. 2003, 2006; Brosius 2003; Kamio et al. 2005; Del Zanna et al. 2006). Recently, Milligan & Dennis (2009) and Chen & Ding (2010) measured Doppler velocities in multi-lines at flare regions to study the evaporation process using Hinode/Extreme-UV Imaging Spectrometer (EIS; Culhane et al. 2007) data. They detected both upflows and downflows in different emission lines.

Chromospheric evaporation can be classified into two types: the explosive evaporation and the gentle one (Fisher et al. 1985a, 1985b, 1985c; Milligan et al. 2006a, 2006b; Brosius 2009). When the lines formed in the upper chromosphere and transition region are redshifted and the hotter lines are blueshifted, this case is considered as an explosive evaporation. When all the lines appear to be blueshifted, it is considered as a gentle one. In hydrodynamic simulations of a flaring atmosphere subject to thick-target electron beam heating, Fisher et al. (1985a) found that an energy flux of about 10¹⁰ erg cm⁻² s⁻¹ can serve as an effective threshold between gentle and explosive evaporations. Note that in observations, explosive evaporation may appear in both large flares and microflares (Milligan et al. 2006a; Brosius 2009; Veronig et al. 2010; Brosius & Holman 2010; Chen & Ding 2010).

Although many observational results support the evaporation model, there are still some inconsistencies between the observations and models. Some hot lines, such as Ca xix, Fe xix, and Fe xxiii, show a dominant stationary component with a relatively weak blueshifted component, indicative of hundreds of km s⁻¹ upflows (Antonucci et al. 1982; Ding et al. 1996; Milligan et al. 2006a; Milligan & Dennis 2009), while the flare dynamic models predict that these lines should be mostly blueshifted. In the case of explosive evaporation, there is a conversion from redshifts to blueshifts. The models (Fisher et al. 1985c) predict that the downflows occur only at temperatures of ⩽1 MK, which was supported by many observations (Kamio et al. 2005; Milligan et al. 2006a; Del Zanna et al. 2006). However, downflows were also detected at much higher temperatures, e.g., 2.0 MK (Milligan 2008). These pose challenges to the theoretical models.

With the high spatial and spectral resolution data of Hinode/EIS, we investigate in detail the process of chromospheric evaporation in a flaring region using multi-lines with different temperatures. Our work focuses on the two problems described above. In addition, we find that opposite line shifts appear in regions of different magnetic polarities. We describe the observations and data reduction in Section 2. The results are presented in Section 3. A discussion is given in Section 4.

The observations presented here are for a GOES C4.2 class flare that started at 02:22 UT on 2007 January 16. The flare is located in the core of NOAA AR 10938. Figure 1 shows the Hinode/EIS Fe xii 195 Å intensity map and the magnetogram measured by the Michelson Doppler Imager (MDI; Scherrer et al. 1995) on board SOHO for the active region (AR). This region presents a bipolar magnetic structure, which is favorable for the production of flares. The 1'' slit of EIS was used to raster over the AR of 240''×240'' with an exposure time of 5 s requiring a total duration of about 26 minutes. EIS scanned this region three times from 01:54:11 UT, corresponding to the preflare, impulsive, and post-impulsive phases of the flare. Our main interest focuses on the second one between 02:20:30 and 02:46:49 UT that covers the impulsive phase. An area with a field of view (FOV) of 80''×50'', marked by the black box in Figure 1, contains the flare and is studied in detail here. Figure 2 shows the Ca ii H images observed by the Solar Optical Telescope (SOT; Tsuneta et al. 2008) on board Hinode, the SOT filter (FG) magnetogram, the Fe xii 195 Å intensity map, and the Doppler velocity of the Fe xii line for this region. The flare has two ribbons as seen from the Ca ii H images. The right ribbon started to brighten at ∼02:30 UT; the left one appeared at ∼02:32 UT and disappeared later than the right one. In particular, we select three points showing significant upflows or downflows in most of the spectral lines under study. From the Ca ii H images and the magnetogram, these points are located at the flare ribbons and close to strong magnetic patches. Point 1 lies in the positive polarity region and at the western part of the right ribbon, while points 2 and 3 lie in the negative polarity region and at the left ribbon. The time for this scan (the second one) is just before the GOES soft X-ray flux reaches its maximum at 02:42 UT. Figure 3 shows the GOES 1–8 Å light curve of the flare and the SOT Ca ii H intensity evolution at the three points in the impulsive phase. We also mark the EIS scanning time ranges and the time when EIS scanned over the three selected points. Note that the GOES light curve shows two peaks, which may correspond to brightenings of the two flare ribbons.

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The EIS spectrum at each location in the raster contains 17 spectral lines. We select 11 of these, spanning a temperature range 0.05–16 MK. Details of the lines are shown in Table 1. The majority of these lines are well resolved with no blends or only trivial blends that can be safely ignored in the AR (Young et al. 2007b). We find that most of the observed line profiles have symmetric Gaussian shapes that can be fitted using a single Gaussian function. However, in some areas, the line profiles are asymmetric and can be well fitted by double Gaussian components.

We reduce the data using the standard EIS software data reduction package. This includes the correction of detector bias and dark current, as well as hot pixels and cosmic ray hits, resulting in absolute intensities in erg cm⁻² s⁻¹ sr⁻¹ Å⁻¹. We also make a correction for a slight tilt of the slit on the CCDs. An additional effect that is corrected for is a variation of line positions over the Hinode orbit due to temperature variations in the spectrometer. Such an orbital variation is obtained by averaging the centroid positions over the length of the slit for which the underlying solar region is mostly a quiet region. We choose the bottom 50 rows in the EIS Fe xii raster (see the red box in Figure 1) as a quiet region. The result is then subtracted from the line center positions measured in all wavelength windows.

EIS does not have an absolute wavelength calibration. We adopt the observed line centers averaged over the quiet region (the red box in Figure 1) as the rest wavelengths. Because of the EIS effective area that only peaks at about 195 and 275 Å in the short-wavelength (SW) and long-wavelength (LW) bands, respectively, and the weak emission in the quiet region, we cannot obtain accurate line centers for some lines, especially the high-temperature lines. Therefore, we use the method of Brown et al. (2007) and the CHIANTI package (Dere et al. 1997, 2009) to determine the reference wavelengths for those lines. To check the reliability of the methods, we use four strong lines, Fe xii, Fe xiii, Fe xiv, and Fe xv, for the test. We find that the wavelengths of the line centers determined using the above methods are nearly the same, with a deviation within ±0.003 Å, which induces a velocity uncertainty of no larger than 5 km s⁻¹.

We co-align the SOT/FG and EIS images by adopting the method described by Guo et al. (2009). Taking into account the instrumental offset between the images taken in the two EIS CCDs, we also shift the LW images by 2'' in the solar X direction and 17'' in the solar Y direction (Young et al. 2007a).

3.1. Point 1: Upflow-dominated in the Positive Polarity Region

EIS scanned point 1 at 02:33:58 UT in the impulsive phase of the flare (see Figure 3). This point shows upflows in all emission lines except for the He ii line. Because of line blending, we do not use the Ca xvii and Fe xxiv lines. We also ignore the Fe xxiii line since it is very weak in this region. We plot the line profiles and their fitting curves in Figure 4. Note that in the wavelength window of the Fe x 184.54 line, there exists the Fe xi 184.41 line; however, the latter is quite distinguishable from the former and therefore does not affect the fitting result (see also Figures 5 and 6). Some of the line profiles are fitted by double Gaussian components. We also calculate the intensity ratio of the blue component to the stationary one. The ratio and the Doppler velocity for the blue component are listed in Table 2. From the results we can see that most of the lines show obvious blueshifts. The upflow velocity increases with temperature from several tens of km s⁻¹ to a maximum of 116 km s⁻¹; the intensity ratio of the two components, derived from the double Gaussian fitting, also has an increasing tendency. For the Fe xiii line, the intensity of the blue component is just 1.72 times that of the stationary component. However, the ratio increases to 9.95 for the Fe xvi line. This indicates that the blue components are dominant over the stationary components for most of the lines, especially those with higher formation temperatures.

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Table 2. Velocity of the Blueshifted Component and Its Intensity Ratio to the Stationary Component at Point 1

Ion	Vd (km s−1)	Ratio
He ii	+5
Fe viii	−28
Fe x	−36
Fe xii	−37
Fe xiii	−68	1.72
Fe xiv	−80	2.53
Fe xv	−89	2.64
Fe xvi	−116	9.95

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We also measure the average Doppler velocity over the nine EIS pixels around point 1 in the preflare and post-impulsive phases and plot this against temperature in Figure 7. The same is done for point 2. Note that the velocities corresponding to the blue components from the double Gaussian fitting are marked with the plus symbols. We find that most of the lines show significant blueshifts or blueshifted components in the impulsive phase, while the shifts are trivial in the preflare and post-impulsive phases.

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3.2. Points 2 and 3: Downflow-dominated in the Negative Polarity Region

We have presented the line profiles and Doppler velocities at flare ribbons of different magnetic polarities to study the chromospheric evaporation process. The key findings of our study are as follows: (1) the flare ribbons with different magnetic polarities can show different patterns of upflows or downflows in the impulsive phase of the flare; (2) the line profiles at one point of positive magnetic polarity are mostly blueshifted, with the blueshifted component dominating over the stationary one especially for hotter lines; (3) downflows are detected at the points of negative magnetic polarity in lines of relatively high temperatures (up to 2.5–5.0 MK); and (4) there exist multi-velocity components in some lines, either cooler or hotter, depending on the points in the flaring regions.

The upflows at point 1 in the positive magnetic polarity region are basically consistent with the scenario of a gentle evaporation process. The most significant finding here is the dominance of blueshifted components in the line profiles, which is consistent with the prediction of theoretical dynamic models. In previous observations, high-temperature lines (e.g., Fe xix, Ca xix, and Fe xxiv) were found to be dominated by a stationary component (Antonucci et al. 1982; Ding et al. 1996; Milligan et al. 2006a; Milligan & Dennis 2009). It is difficult to explain such observational results in terms of the basic dynamic models. Doschek & Warren (2005) argued that the stationary component is from the top of the flare loop where the evaporated mass is accumulated, or that the mass is moving perpendicular to the line of sight. They also suggested the possibility that the instrumentation is not sensitive enough to detect the earliest blueshifted emission; by the time the emission level has risen sufficiently, the flare loops have already been filled. Falewicz et al. (2009) took into account the geometrical dependence of the line-of-sight velocities of the plasma motions along the loops inclined toward the solar surface, as well as a distribution of the flare sites over the solar disk. They concluded that the stationary component can be observed for all flares during their early phases of evolution; on the other hand, the blueshifted component may be undetectable even for plasma moving along the flaring loop with a very high velocity. In spite of these explanations, searching for a dominant blueshifted component in high-temperature lines during flares is still an important task for reducing the discrepancy between observations and models. In this work, although we cannot accurately determine the Doppler velocities of the Fe xxiii (12.5 MK) and Fe xxiv (15 MK) lines at point 1 because of the EIS sensitivity limitation and line blending, we confirm that some line profiles are dominated by blueshifted components. Our result further shows that the intensity ratio of the blueshifted component to the stationary one increases with temperature.

The counterpart of upflows is downflows in the flaring loop subject to a rapid energy deposition and momentum balance. Downflows were mostly detected in chromospheric lines before the SOHO era (e.g., Ichimoto & Kurokawa 1984; Ding et al. 1995). This is known as the phenomenon of chromospheric condensation (Fisher 1989). With space instruments, downflows are also seen in some coronal lines. Most previous observations showed that the downflows occur at temperatures of ⩽1 MK (Kamio et al. 2005; Milligan et al. 2006a; Del Zanna et al. 2006). Quite recently, Milligan & Dennis (2009) detected downflows at a temperature of 1.5 MK. Moreover, Milligan (2008) reported a redshift of 14 km s⁻¹ in the Fe xv line (2 MK) at flare footpoints. The fact that the downflows exist at such a high temperature provides a constraint to the dynamic model. Our measurements presented here show that redshifts can appear in lines of even higher temperature (Fe xvi, Ca xvii) at the points in the negative magnetic polarity region. Most of the line profiles at point 2 show redshifts, while only very few are blueshifted. This case can be regarded as an explosive evaporation (Milligan et al. 2006a; Brosius 2009). If we take the division temperature between upflows and downflows as the site of energy deposition by the electron beam, the appearance of the unusually high division temperature at point 2 implies a rather high energy deposition site. A possible reason is that the flare loop has been filled with enough mass before the impulsive phase. A high coronal density can prevent the electron beam from penetrating deeper. Here we do not detect blueshifts at temperatures higher than 15 MK due to the limitation of the EIS dynamic range, but we find that the Fe xxiii and Fe xxiv line intensities at this point are still strong and their profiles are well Gaussian-shaped in this event. This seems to support the hypothesis of a high coronal density. In addition, Brosius (2003) detected redshifted emission that persisted at least 20 minutes after the cessation of blueshifts and suggested that flare plasma, heated and accelerated upward during the impulsive phase, subsequently cooled and fell back down in what may be thought of as "warm rain." In our investigation, scanning of the AR by EIS lasted about 26 minutes; therefore, we do not have a high enough temporal resolution here. However, from Figure 7, we find that the downflows still exist in the post-impulsive phase at point 2, contrary to the blueshifts at point 1, which diminish to nearly zero in the post-impulsive phase. From the SOT Ca ii H movie, we find that the flare ribbons are not brightened simultaneously. Therefore, we cannot exclude the warm rain as a possible cause of the downflows.

The three points discussed here are located in different magnetic polarity regions. They show different patterns of mass flows. Point 1 lies in the positive polarity region, while points 2 and 3 lie in the negative polarity region. We further check the magnetic connectivities between the positive and negative magnetic polarities. The available data cannot provide a definite conclusion regarding whether point 1 is magnetically connected to point 2 or 3. However, even if they belong to different magnetic loops, their different evaporation patterns shown suggest that in one flaring region, the heating mechanisms and atmospheric conditions may vary from point to point.

The authors thank the referee for valuable comments on the paper and Feng Chen, Xin Cheng, and Zongjun Ning for helpful discussions. This work was supported by NSFC under grants 10828306 and 10933003 and by NKBRSF under grant 2011CB811402. Hinode is a Japanese mission developed and launched by ISAS/JAXA, collaborating with NAOJ as a domestic partner, and NASA (USA) and STFC (UK) as international partners. Scientific operation of the Hinode mission is conducted by the Hinode science team organized at ISAS/JAXA. Support for the post-launch operation is provided by JAXA and NAOJ (Japan), STFC (UK), NASA, ESA, and NSC (Norway).