Height Dependence of Gas Flows in an Ellerman Bomb

Transcription

1 PASJ: Publ. Astron. Soc. Japan 60, , 2008 February 25 c Astronomical Society of Japan. Height Dependence of Gas Flows in an Ellerman Bomb Takuma MATSUMOTO, 1,2 Reizaburo KITAI, 2 Kazunari SHIBATA, 2 Kenichi OTSUJI, 1,2 Takuya NARUSE, 1 Daikou SHIOTA, 2 and Hiroyuki TAKASAKI 2,3 1 Department of Astronomy, Kyoto University, Sakyo-ku, Kyoto Kwasan and Hida Observatories, Kyoto University, Yamashina-ku, Kyoto Accenture Japan Ltd., Akasaka Inter City, Akasaka Minato-ku, Tokyo mtakuma@kwasan.kyoto-u.ac.jp (Received 2006 November 25; accepted 2007 September 20) Abstract We performed spectroscopic observations of Ellerman bombs (EBs) in an active region of NOAA at Hida Observatory on 2004 November 24. The photospheric velocity fields of EBs have for the first time been investigated spectroscopically. Fromthe Doppler shifts ofa Ti II absorption line ( Å) and a broad H emission line, we derived the photospheric velocity and the lower chromospheric velocity, respectively. The photospheric velocity of EBs was 0:2 km s 1, indicating downward flow, on average. We found that the photospheric velocity variation of EBs has a good temporal correlation with the H wing emission variation. On the other hand, the chromospheric velocity showed an upward flow of 1 3 km s 1 on the average. From the characteristics of the flow field, we conclude that the observed EB occurred at the upper photospheric level. We suggest that it is important to know the motions of EBs in the photosphere because a plausible triggering mechanism of EBs is magnetic reconnection in the low-lying atmosphere. Key words: Sun: activity Sun: atmospheric motions Sun: chromosphere Sun: magnetic fields Sun: photosphere 1. Introduction Ellerman bombs (EBs: Ellerman 1917), or moustaches, are interesting explosive events in the solar chromosphere, which show characteristic H spectra, broad emissions with central absorptions. The typical size of EBs is on the order of 1 00,and their lifetimes are generally estimated to be min. They generally occur in emerging flux regions, beneath expanding arch filament systems. EBs appear as elongated structures, and are occasionally associated with small surges. Although many previous works have shown the basic properties of EBs (Roy 1973; Kurokawa et al. 1982; Kitai 1983; Kitai & Muller 1984; Zachariadis et al. 1987; Nindos & Zirin 1998; Georgoulis et al. 2002; Pariat et al. 2004, 2007), their triggering mechanism has not yet been clearly understood. The H line-broadening mechanism has been studied by many authors. Typically, the Doppler widths of EBs emission at the H line wing are 3Å, and correspond to Doppler velocities of 150 km s 1. Because these velocities greatly surpass the chromospheric sound speed (10 km s 1 ), it is hard to consider that microturbulent motions produce such broad emission components. On the other hand, the indices of powerlaw functions fitting the wing profile prevent the simple Stark effects from explaining the line-broadening mechanism, since they do not produce power-law profiles, but damping profiles. However, Kitai (1983) showed that non-lte calculations of a perturbed, heated and condensed chromosphere well reproduced the observed H profiles. Hénoux, Fang, and Ding (1998) also reproduced profiles of EBs by considering energetic particle beams bombarding the atmosphere. Both studies suggested that heating of the lower atmospheres, such as the lower chromosphere, plays an important role in producing the H profiles of EBs. A plausible triggering mechanism of EBs is magnetic reconnection occuring in the low-lying atmosphere. EBs are often observed under arch filament systems where magnetic flux is expected to be in a form of undulatory tubes at photospheric levels (Pariat et al. 2004). Almost half of the EBs are seen to occur at magnetic dips of the emerging flux (Pariat et al. 2004). As the flux tubes emerge, magnetic field lines with opposite directions approach, thus reconnecting at the magnetic dips. As a result, we can expect heating and bidirectional flow at the low-lying atmosphere. Although magnetic reconnection is considered to be the mechanism of primary energy release of solar flares at coronal heights, it may occur in lower layers, and result in EBs. Recently, Chae, Moon, and Park (2003) proposed that magnetic reconnection occurs at different levels of the photosphere and the chromosphere without any preferred height. There have been various observations on the gas velocity in EBs at different height layers. At the photosphere, Severny (1968) reported 1 3 km s 1 upward flows in EBs atmospheres. On the other hand, Georgoulis et al. (2002) conducted statistical studies on EBs, and found that more than 80% of EBs are associated with km s 1 downward flows. At the chromosphere, 6 km s 1 upward flows have been observed on several EBs (Kitai 1983). Kurokawa et al. (1982) also found 8kms 1 upward flows during the brightening phase of EBs. Chen, Fang, and Ding (2001) reproduced these upflows by numerical simulations of magnetic reconnection in the chromospheric layer. Roy and Leparskas (1973) have observed dark elongated matter ejected from EBs to coronal heights with

2 96 T. Matsumoto et al. [Vol. 60, Fig. 1. Active region NOAA observed on 2004 November 24. The upper left and the upper right panels are, respectively, H 0.8 Å filtergrams observed at HIDA observatory. The lower left panel is the SOHO/MDI magnetogram. The lower right panel is the continuum image observed at HIDA. HL indicates the locations of hair lines. An Ellerman bomb is on the slit. a 100 km s 1 upward velocity. To obtain clues for solving the triggering mechanism, it is necessary to know the height dependence of the EB structures. The location of the primary energy input is key information. A good correlation is reported between the H wing emission of EBs and TRACE 1600 Å bright points, suggesting that heating in the photosphere, in the temperature minimum region, is very important for producing EBs (Qiu et al. 2000; Georgoulis et al. 2002). If local atmospheric heating is due to magnetic reconnection, bidirectional flows are expected in EBs. In this study, we thus concentrated on obtaining the height dependence of gas flows by a spectroscopic method. In order to achieve this purpose, we conducted a high-dispersion time-sequenced spectroscopic study in H and a photospheric line. The present observational study on the height dependence of gas flows in EBs has shown, as an important result, that the energy input site lies between the lower photosphere and the lower chromosphere. 2. Observation We succeeded to obtain good H spectrograms and H filtergrams of EBs in our observation on 2004 November 24 at Hida Observatory. There were many EBs in the active region of

3 No.1] GasFlowsinanEllermanBomb 97 Fig. 2. H spectrogram of the EB. The broad emission line and the central absorption line of EBs can be seen at the center. Fig. 4. Procedure to define the lower chromospheric velocity. We determined the bisectors of the H emission wing at 20 intensity levels. The average wavelength of these points is assumed to be the center of the H emission component. Fig. 3. Time evolution of the H excess emission profiles of the EB. Both the gray color and the contour represent the intensity integrated around the H wing. There are two EBs in our observation (EBa, EBb). The chromospheric velocity was derived by measuring the Doppler shift of these wing profiles. NOAA 10705, located at W70 ı S5 ı near the west limb. NOAA had a preceding sunspot and a following plage, and is hence called p type. There was also an arch filament system in this region (figure 1). The purpose of our observation was to search for the height dependence of gas flows in EBs, and to find the energy-release sites of EBs. We set the spectrographic camera so that its FOV covered both the H emission line and the Ti II absorp- tion line ( Å) simultaneously in one frame, which permitted us to do a multi-wavelength observation. Moreover, in order to estimate 5-min oscillations in the photosphere, and to increase the measurement accuracy of the photospheric velocity, we made a long time-series observation of about 60 min (00:36:37 UT 01:36:48 UT). Grating spectrograms were obtained by the vertical spectrograph of the Domeless Solar Telescope (DST) at Hida observatory (figure 2). The spectrographic field of view was Å, with a 1000 pixel 1000 pixel format and a cadence of spectrogram acquisition of 5 s. H filtergrams were obtained by the Lyot filter of DST and taken at the H center, H 0.5 Å,H 0.8 Å, and the continuum. The spatial resolution of H filtergrams was and the cadence of filtergram acquisition was 6 images per 30 s. The FOV of our filtergrams was , taken in the 2000 pixel 2000 pixel format. 3. Data Reduction Dark-current subtraction and flat-fielding were applied in a standard manner for all of the spectrographic images. A wavelength calibration was performed with reference to four terrestrial H 2 O absorption lines ( Å, Å, Å, Å). A Gaussian fitting for these

4 98 T. Matsumoto et al. [Vol. 60, Fig. 5. Time evolution and spatial distribution of the 5-min oscillation velocity at the photosphere. The gray color represents the magnitude of the velocity. The amplitude of the 5-min oscillation velocity is 0.1 km s 1 on average. The lower image corresponds to the H A filtergram and the thick solid line in the image indicates the location of the slit. absorption line profiles permitted us to identify the position of the line centers with an accuracy of A, corresponding to 0.1 km s 1 at H, where the wavelength dispersion of the spectrogram was 0.01 A pixel 1. A wavelength calibration was made by cubic interpolation. H lines were analyzed after deriving enhanced H emission profiles of EBs, which were obtained from the observed profiles by subtracting neighboring unperturbed H absorption profiles. The time evolution of the intensity excess profile of the observed EB is shown in figure 3. We can see characteristic H profiles of EBs, and that the wide emission profiles were depressed by the central absorption component. The Doppler shifts of the H wide emission component were derived from the bisector of the emission components. See figure 4 for details of the method. A typical value of the measured velocity was 2 4 km s 1. The obtained Doppler velocity includes, beside the EB velocity, the steady components such as solar rotation, Earth orbital motion etc, and also the temporally variable component of 3-min oscillation. By taking the central wavelength of the unperturbed H absorption as a reference, we removed the steady components with 1 km s 1. Because the velocity amplitude of 3-min oscillation near the location of EBs was 0.2 km s 1, this component was small enough to be neglected. The photospheric velocity was derived from the Doppler Fig. 6. Time evolution of the EB during 40 min from the start time of our observation. The upper panel shows the intensity variation of the H wing. The middle panel shows the photospheric velocity obtained from the Doppler shift of the Ti II absorption line. The bottom panel shows the chromospheric velocity obtained from the Doppler shift of the H wing emission component. For all of the plots, cross marks represent the observed values and the smooth lines show slowly varying Fourier components of the observed values. shift of the Ti II absorption line ( A ), whose formation height is considered to be in the photosphere. There are also two velocity components in the derived photospheric velocity besides the EB velocity: the steady components and the temporally variable component of the 5-min oscillation. In order to remove the steady component of the velocity, the spatially averaged velocity of the surrounding quiet region was subtracted.

5 No.1] GasFlowsinanEllermanBomb 99 Fig. 7. Time evolution and spatial distribution of the photospheric velocity. The upper panel shows the photospheric velocity. Dark regions represent blue shifted regions. White contours represent the temporal and spatial distribution of the integrated emission of the EB. The bottom image shows the H Å filtergram and the thick solid line in the image indicates the location of the slit. The 5-min oscillation component was subtracted by Fourier analysis. Time profiles and the spatial distribution of the 5-min oscillation component are shown in figure Results and Discussion From our spectrograms we obtained time profiles of the H wing intensity, the photospheric velocity, and the chromospheric velocity of the EB. From the H filtergrams, we identified a surge associated with the EB at the end of its lifetime. Figure 6a shows the time evolution of the H wing intensity of the observed EBs. The cross marks in the figure represent the observed values, and the smooth line shows slowly varying components, whose Fourier period is longer than 6 min. Smoothing in time was introduced to suppress the 5-min oscillation component and high-frequency noise. The wing intensity was obtained by integrating the H wing from 3 Åto 2Å and from 2 Åto3ÅaroundtheH center. The intensity was then spatially averaged over 1 00 and normalized by the continuum value of H. TheH intensity time variation has two significant peaks at 00:49:26 UT and 01:18:52 UT, where the noise level is 1%. These peaks are considered to correspond to two distinct and adjacent EBs (EBa, EBb); otherwise, the lifetime of about 30 min would be too long for a single EB. We confirmed from H filtergrams that EBa was on the Fig. 8. Small surge in the EBb. The above four images correspond to the time series of H Å filtergrams. The central bright point is the EBb and the arrows show the location of the surge ejecta. White contours in the above panels are the intensity levels for the dark region. The thick solid line in the images indicates the location of the slit. slit position during our observation, so we could obtain the Doppler shifts of the gas flows for the EBa. On the other hand, EBb was out of the slit position from 01:20:00 UT,atthelast phase of our observation. Although spectroscopic results for the EBb could not be obtained, the direction of the gas flows of a small surge associated with the EBb could be derived from a time-sliced analysis of H images. Figure 7 shows the time evolution of the photospheric velocity field along the slit. Gray scale colors represent the photospheric velocity, such that dark regions represent blueshifted regions. White contours are 1% and 2% of the excess intensity, which represent the location of the EB. A close correlation exists between the photospheric blue-shifted region and the region of the EB. The spatially averaged time profiles of the photospheric velocity at the position of the EB are shown in figure 6b. The cross marks and the smooth line have the same meanings as in figure 6a. The photospheric velocities of the EBs indicated km s 1 blue shifts during its lifetime. At the end of the lifetime of the EB, the photospheric velocity became zero. We also obtained the chromospheric velocity from the Doppler shift of the H emission component. The time evolution of the chromospheric velocity is shown in figure 6c. The EBs had red-shift velocities of 1 3 km s 1 at the chromosphere. Note that the velocity from 00:52 to 01:13 is not reliable, since there is a too-low excess emission to obtain the Doppler shift of the emission component. We observed a small surge in the EBb (figure 8). A timesliced image of the surge is shown in the right part of figure 9.

6 100 T. Matsumoto et al. [Vol. 60, Fig. 9. Time-sliced image of a small surge associated with the observed EBb. The left panel is the H Å filtergram and the thick solid line in the image indicates the location of the slit. The right panel is a time-sliced image using the intensity of H Å along the fixed measuring line (white solid line in the left image). The white dotted line shows the path of the surge. The speed of the surge is nearly 5 km s 1 at the upper chromosphere. Fig. 10. Possible trajectory of the surge from EBb. The co-latitudinal plane seen from the North is shown. The plane is separated into regions I to IV by Cartesian axes located at the EBb. The origin represents the reconnection point. The axis of the reconnection flow in EBb is considered to lie on regions I and III. Fig. 11. Schematic image of the inclined EB configuration observed near the limb. Reconnection jets are ejected in the upward and downward directions. In this configuration, downward flows are observed as blue shifts and upward flows as red shifts. At the time of ejection, the EB was out of the slit position, so we used an intensity of H Å along a fixed measuring line (white line in the left panel of figure 9). In the right panel of figure 9, EBb is located at the lower region, and the dark absorption matter is ejected along the white dotted line. The ejection velocity of the surge is found to be 5 km s 1. Since our target region was located at 70 degrees west, projection effects must be considered to obtain a correct view of the event. First, we assume the triggering mechanism of EBs to be magnetic reconnection. Gas flows in EBs should then have a jet-like structure. Along the jet axis, bidirectional flows are expected to occur from the reconnection point. In figure 10, we sketch the possible geometrical configuration of the jet. The origin represents the reconnection point where the surge from EBb is ejected. Since the surge from EBb moved towards the west limb in our filtergram observation, the surge must be ejected to region I or II of figure 10. Moreover, because the surge was observed with H Å, the direction of the surge ejection is limited to region I. Hence, the axis of bidirectional flow must be inclined towards the limb by more than 20 degrees. In addition, if the surge moved downward to the dense solar atmosphere, it would suffer significant deceleration. Figure 9 shows that the surge ejecta moved nearly 3000 km without any significant deceleration. Thus, finally, we think that the surge was ejected to region I, and to higher levels. EBa and EBb occur at nearly the same location, and follow each other by less than 30 min. We assume that the magnetic configuration does not change very much between EBa and EBb. Hence, the axis of EBa must lie along the same direction as that of EBb. A photospheric blue shift therefore represents a downward flow, while a chromospheric red shift represents an upward flow.

7 No.1] GasFlowsinanEllermanBomb 101 The schematic configuration that we considered is shown in figure 11. Taking projection effects into consideration, we can interpret the above observation as a 0.2=sin km s 1 downward flow at the lower photosphere, 1 3=sin km s 1 upward flow at the lower chromosphere, and 5=sin km s 1 at coronal heights, where is the angle between the sky plane and the jet direction. Thus in this case, the reconnection site probably lies between the lower photosphere and the lower chromosphere. In order to investigate the detailed height dependence of the gas flow in EBs, it is necessary to know the formation heights of the Ti II absorption line and the H emission line. We define the photosphere as the layer where the optical depth becomes unity ( = 1). The altitude of the photosphere is set to 0 km. The formation height of Ti II is considered to be in the lower photosphere from the velocity amplitude of the 5-min oscillation. The maximum velocity amplitude of the 5-min oscillationwasobservedtobe0.1 km s 1 at W70 (figure 5). Due to the fact that the 5-min oscillation is a vertical mode, the real root-mean-square of the observed amplitude is considered to be 0.2 km s 1. This amplitude corresponds to the lower photospheric oscillation amplitude (Canfield 1976) (0 200 km). The formation height of the H emission line is considered to be the lower chromosphere ( km) from non-lte calculations (Kitai 1983). Hence, the observed bidirectional flows probably start between the lower photosphere and the lower chromosphere ( km), and the reconnection site also lies between these layers. Let us discuss our observational results on EBs from the viewpoint of magnetic reconnection triggering. EBs are often associated with small surges (Roy 1973), likely to expand and shrink with an elongated shape (Kurokawa et al. 1982). Moreover, the spatial scale of EBs is larger in the chromosphere than in the photosphere (Kitai & Muller 1984; Zachariadis et al. 1987). Therefore, EBs are considered to have funnel-like structures. These observations imply that EBs eject bidirectional jets from the reconnection site. The directions of funnel-like structures of EBs would depend on the magnetic field configuration. The major theoretical scenario for EBs is the magnetic reconnection scenario (a working hypothesis). Subphotospheric magnetic flux may rise by magnetic buoyancy, and then form current sheets at the magnetic dips, in which a type of resistive instability can take place (Pariat et al. 2004). The purpose of this paper is to present evidence supporting the magnetic reconnection scenario for EBs. Assuming this scenario, upward flow and downward flow will occur above and below the reconnection point. We consider that the bidirectional flow found here at the low atmosphere corresponds to reconnection flow. Moreover, the observed flow velocity of a few km s 1 is nearly the same as the theoretically predicted reconnection flow velocity at the photosphere (Chae et al. 2003). These results support the magnetic reconnection scenario for EBs. The spectral characteristics of EBs, a broad H emission wing and a central absorption core, are well explained by local heating in the low atmosphere (Kitai 1983; Ding et al. 1998; Hénoux et al. 1998). Magnetic reconnection in the low atmosphere with weakly ionized plasma is possible, and consists of a plausible triggering mechanism of EBs. The effects of neutral atoms on the resistive instability produce a typical lifetime of EBs of 10 min (Li et al. 1997). Takeuchi and Shibata (2001) investigated photospheric magnetic reconnection by numerical simulations, and obtained bidirectional flows. Our spectral observation for gas flows in an EB shows that bidirectional flow starts between the lower photosphere and the lower chromosphere. We thus suggest that magnetic reconnection can indeed occur in such a layer, and give rise to heating in the low atmosphere. The result that EBs are low-atmospheric events agrees with the result of Qiu et al. (2000) and Georgoulis et al. (2002). Also, we have found new evidence for low-altitude reconnection in EBs from our multi-line spectroscopic analysis. 5. Conclusion Although EBs are suggested to be the result of magnetic reconnection in the lower atmosphere (Kitai 1983; Ding et al. 1998; Hénoux et al. 1998; Qiu et al. 2000; Chen et al. 2001; Georgoulis et al. 2002; Chae et al. 2003; Pariat et al. 2004; Fang et al. 2006; Socas-Navarro et al. 2006), there has been few spectroscopic observations about the height dependence of the gas flows in EBs. In order to study the gas-flow patterns in EBs, and to confirm lower atmospheric reconnection, we conducted spectroscopic observations of two EBs as a case study, and obtained spatial distributions and the time evolution of the intensity, photospheric velocity, and chromospheric velocity of the EBs. For EBa, we found a km s 1 downward flow in the lower photosphere, and a 1 3 km s 1 upward flow in the lower chromosphere. For the second EB, which is associated with a small surge, we found that an EB can have a 5 km s 1 upward flow velocity in the upper chromosphere. Therefore, we concluded that the reconnection site of the EBs was between the lower photosphere (0 200 km) and the lower chromosphere ( km), where we define the photosphere as the = 1 layer with an altitude of 0 km. Although our result supports the low-altitude reconnection scenario for EBs, further work is required to unambiguously prove or reject the plausible reconnection mechanism. We must observe many other EBs in the same way at various locations on the disk, with various Fraunhofer lines. The time evolution of the magnetic field would also give interesting information. Numerical simulations of emerging flux, whose reconnection point is in the lower atmosphere, will help us to finally identify the triggering mechanism of EBs. The authors are grateful to all staff members of the Kwasan and Hida Observatories, Kyoto University for their guidance in our spectroscopic observations with the Domeless Solar Telescope and for the fruitful discussions with them. We also thank the instrumental teams of SOHO missions for their opendata policies. The authors are supported by a Grant-in-Aid for the 21st Century COE Center for Diversity and Universality in Physics from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan. This work was supported by the Grant-in-Aid for Creative Scientific Research The Basic Study of Space Weather Prediction (17GS0208, Head Investigator: K. Shibata) from the Ministry of Education, Culture, Sports, Science and Technology of Japan.

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