A Proposal Summary

The relationship between coronal mass ejections (CMEs) and filament eruptions has been among most basic topics in solar physics. The study of it in the context of magnetic structure and evolution is even more important to understand the triggering mechanisms of CMEs. However, a number of studies including ours have been inconclusive. The major difficulty is that the CMEs is most visible as limb events, while filament eruption and magnetic topology evolution can be best studied when a source region is on the disk. The unprecedented data from Earth Connection Coronal and Heliospheric Investigation instrument (SECCHI) on board STEREO provide unique opportunity to revisit this problem, as CMEs and associated eruptions will be viewed at different angles from earth-based observatories. Therefore, both CME dynamics and disk signature can be studied simultaneously. We propose to perform three related tasks.

(1) In next three years, we will study all available events of filament disappearance observed by our Global High Resolution Halpha Network and study their response as observed by white-light coronagraphs and EUV imager of SECCHI. We anticipate that over 100 events will be studied and eruption of smaller filaments may have associated CMEs as well. We will construct eruption trajectories from chromosphere through corona. In addition, we will determine the differences between eruptions associated with CMEs and those that are not associated. The determining factors may include the structure (e.g. chirality, and furthermore, the helicity) of the overlying field structure, the size and the shape (e.g. sigmoid or not) of the filaments, whether flare associated or not, whether the filament reformed, and whether it is an active region or quiet-sun filament.

(2) We will select 20 to 40 events from the above filament list, to calculate one of the two components of helicity injection: helicity transport into the solar atmosphere by photospheric footpoint motions (shuffling term). We realize the difficulty in studying the the second term: the advection term due to flux emergence. So we will make sure that for the selected events, the advection term can be neglected. The selected events will be close to disk center (earth-based observation). This helicity calculation will be first based on SOHO MDI data and to be compared with BBSO high resolution magnetograms when available. The transition will be made to use HMI data when Solar Dynamic Observatory is launched. We will study the evolution of helicity injection and compare with the timing of the onset of filament eruptions and CMEs, and correlate the amount of helicity injection and CME kinetics.

(3) We will explore the correlation between the kinetic properties of CMEs and the accompanying magnetic reconnection rate, i.e, the rate of magnetic flux change involved in magnetic reconnection in the low corona. We will use Halpha images obtained from our our global Halpha network as the primary data source to trace the flare ribbon expansion and the filament rising motion. The magnetic reconnection rate is deduced by measuring the expansion of flare ribbons across the magnetic field. We anticipate that this research will provide a clear evidence of the important role of magnetic reconnection in driving large scale eruptions that travel through the interplanetary space. With our limited sample studied so far, a positive correlation is found between the magnetic reconnection rate and the acceleration of erupting filaments that represent the early stages of large-scale solar eruptions. Even a better correlation was found between the total reconnected magnetic fluxes and the CME speeds. We propose to continue and advance this study with the help of newly developed program to automatically detect and track flare ribbons and CME fronts, so we can monitor the magnetic reconnection rate and predict the resulting CME speed in real time.

B Introduction and Objectives of Proposed Research

Filament eruptions, flares, and CMEs are the most important solar events as far as space weather effects are concerned linking solar eruptions, major interplanetary disturbances and geomagnetic storms (Gosling et al., 1991). A halo CME, which is usually associated with activity near solar disk center, has great influence on space weather because of the likelihood to launch geoeffective disturbances toward Earth Cane et al., 2000; Webb et al., 2000. The sign of magnetic helicity in active regions can be used to predict the orientation of the magnetic field associated with a CME and furthermore, the probability of a geomagnetic storm (Yurchyshyn et al., 2000).

In order to gain a better understanding of CMEs and to improve the reliability of geomagnetic storm predictions and warnings, it is essential to observe early manifestations of CMEs in the solar atmosphere. Thus,our goal is to find possible relationships between solar surface phenomena such as filament eruptions and flares, the CMEs occurrence, and the properties of the associated magnetic field. Once a relationship is found, it can serve as an indicator for the occurrence of geomagnetic storms.

Filaments and prominences refer to the same physical structures on the Sun, either projected onto the disk or extending above the limb. The majority of previous studies have focused on prominences because they could easily be detected, observed, and measured against the dark sky background. Moreover, CMEs, associated with the prominences, are not difficult to detect. Many prominence classifications have been proposed in the past. For example, (Gilbert et al., 2000) developed definitions of active prominences (APs) and eruptive prominences (EPs) and studied the relationship between APs, EPs, and CMEs for 54 events. They found that 94% of the EPs had an associated CME compared to only 46% for APs. Gopalswamy et al. (2003) defined a prominence as a radial or a transverse event. Authors showed that the radial events have a strong correlation to the CMEs: 83% of the radial events were associated with CMEs compared to 24% for transverse events. However, Yang and Wang (2002) showed that the connection between filament disappearances observed in Hα spectral line and CMEs is weak ranging from 10% to 30%, depending on heliographic location of the source regions. Based on similar data base, but with larger filament threshold, the percentage of filament eruption associated with CMEs is increase drastically (Jing et al., 2005).

Filament disappearance does not always imply filament eruptions. Depending on their physical nature, disappearing filaments can reappear. Two processes have been proposed by (Mouradian et al., 1995): dynamic sudden disappearance (DSD) and thermal disappearance (THD). DSD is due to restructuring of the magnetic field and it ultimately leads to the disappearance of the filament, whereas THD is due to heating of the plasma in the filament, which will reappear once it cools down. Since THD is not related to magnetic field reorganization, we excluded it from this study.

In the last few years, Big Bear Solar Observatory (BBSO) developed its new generation of full-disk Hα observations. Every minute, a full-disk Hα image is recorded by a 2048×2048 Kodak Megaplus camera with the 8-inch Singer-Link telescope. The pixel size of the images is 1′′, which yields a spatial resolution of 2′′. Because of BBSO’s long-lasting clear skies and good seeing conditions, these movies provide a unique tool to monitor filaments continuously during their disk passage. At the same time, BBSO provides high resolution Hα data with its 65cm and 25cm telescopes. These observations cover only a fraction of solar disk, but images have a resolution around 0.5 arcsec and cadence about 15 s.

Based on our successful experience in making high resolution full disk Hα observations at BBSO, and the urgent need to monitor the evolution of the solar chromosphere continuously, we have deployed a network for Hα full-disk observations utilizing existing facilities at BBSO in California, the Kanzelhohe Solar Observatory (KSO) in Austria, Catania Astrophysical Observatory (CAO), Yunnan Observatory (YO), and Huairou Solar Observing Station (HSOS) in China. All these observatories have over 200 sunny days a year, good seeing conditions, a sufficient and experienced observing and research staff and well-established Hα telescope systems. Each observatory is now equipped with a 2K×2K digital camera (except CAO which uses 1K by 1K camera).

The study of CMEs and filament eruptions in the context of magnetic structure and evolution is important to understand the triggering mechanisms of CMEs. However, a number of studies including ours have been inconclusive. The major difficulty is that the CMEs is most visible as limb events, while filament eruption and magnetic topology evolution can be best studied when a source region is on the disk. The unprecedented data from Earth Connection Coronal and Heliospheric Investigation instrument (SECCHI) on board STEREO provide unique opportunity to revisit this problem, as CMEs and associated eruptions will be viewed at different angles from earth-based observatories. Therefore, both CME dynamics and disk signature can be studied simultaneously. Coordinated observation provide unique opportunity to reveal 3-D magnetic structure of eruptive events from filament eruption to CME. Our proposed research will be targeted at the following areas:

(1) In next three years, we will study all available events of filament disappearance observed by our Global High Resolution Halpha Network and study their response as observed by white-light coronagraphs and EUV imager of SECCHI. We anticipate that over 100 events will be studied and eruption of smaller filaments may have associated CMEs as well. We will construct eruption trajectories from chromosphere through corona. In addition, we will determine the differences between eruptions associated with CMEs and those that are not associated. The determining factors may include the structure (e.g. chirality, and furthermore, the helicity) of the overlying field structure, the size and the shape (e.g. sigmoid or not) of the filaments, whether flare associated or not, whether the filament reformed, and whether it is an active region or quiet-sun filament.

(2) We will select 20 to 40 events from the above filament list, to calculate one of the two components of helicity injection: helicity transport into the solar atmosphere by photospheric footpoint motions (shuffling term). We realize the difficulty in studying the the second term: the advection term due to flux emergence. So we will make sure that for the selected events, the advection term can be neglected. The selected events will be close to disk center (earth-based observation). This helicity calculation will be first based on SOHO MDI data and to be compared with BBSO high resolution magnetograms when available. The transition will be made to use HMI data when Solar Dynamic Observatory is launched. We will study the evolution of helicity injection and compare with the timing of the onset of filament eruptions and CMEs, and correlate the amount of helicity injection and CME kinetics.

(3) We will explore the correlation between the kinetic properties of CMEs and the accompanying magnetic reconnection rate, i.e, the rate of magnetic flux change involved in magnetic reconnection in the low corona. We will use Halpha images obtained from our our global Halpha network as the primary data source to trace the flare ribbon expansion and the filament rising motion. The magnetic reconnection rate is deduced by measuring the expansion of flare ribbons across the magnetic field. We anticipate that this research will provide a clear evidence of the important role of magnetic reconnection in driving large scale eruptions that travel through the interplanetary space. With our limited sample studied so far, a positive correlation is found between the magnetic reconnection rate and the acceleration of erupting filaments that represent the early stages of large-scale solar eruptions. Even a better correlation was found between the total reconnected magnetic fluxes and the CME speeds. We propose to continue and advance this study with the help of newly developed program to automatically detect and track flare ribbons and CME fronts, so we can monitor the magnetic reconnection rate and predict the resulting CME speed in real time.

C Research Plan

C.1 Relationship Between Filament Disappearances and CMEs

As we mentioned, Yang et al. (2001) studied about 500 filament disappearances observed in Hα at BBSO and compared them with the corresponding LASCO data. Filament data are from the 2K×2K BBSO Hα archive. Figure 1 shows the number of filament disappearances as a function of the distance to the solar Figure 1: Distribution of filament disappearance as function of the distance from the disk center. The lighter

histogram is the total number of events, and the darker histogram is the number of events with associated CMEs found from the LASCO data. We also indicated at the top of each bar the percentage of events for which LASCO events have been identified.

disk center. The normalized distance mean the physical distance to the solar disk center divided by the solar radius. The lighter histogram is the total number of events, and the darker histogram is the number of the events for which we can find associated CMEs from the LASCO data. We also marked the percentage of events having an associated CME. Different from other studies, even limb events do not have a high percentage of associated CMEs. Please note that we have already considered the possibility that some filaments may have disappeared due to heating, not eruption. Based on our high cadence movies, we found that thermal filament disappearance is a very small (10% or less) fraction.

One of the notable examples is the October 13, 1999 event, which is shown in Figures 2 and 3. Figure 2 shows a sequence of BBSO full-disk Hα images from October 11 to 15, 1999. The filament eruption occurred at 07UT on October 13. Two magnified images are included to demonstrate the high quality of the full disk data. Figure 3 shows the corresponding CME. This event clearly fits the classical CME picture with a bright rim-like arcade structure, a cavity, plus a core of filament material (Hundhausen, 1988).

As we mentioned, combining SECCHI Coronagraph data with continuous Hα data will provide powerful tool to obtain meaningful statistics between filament eruptions and CMEs

Description of SECCHI instrument, include a picture.

The study of new events will be based on the observations of the Global Hα Network. Five stations of our Hα network routinely record full-disk Hα data with a cadence of one minute using a 2K×2K Kodak Megaplus camera. The images are taken from sunrise to sunset. Roughly 1,500 to 2,000 images are obtained each day. At the end of each observing day, a quick video replay of the observations is reviewed by one of the observers in each station, and any notable filament activity is recorded. The Hα filter is tuned 3 to 4 Figure 2: The sequence of Hα full-disk images demonstrating the disappearance of a major filament inside

the white box. A Kitt Peak magnetogram at 15:43 on October 13 is also included to show the magnetic configuration. In the second row, we show a magnified field of view from the same images in the first row to demonstrate that the fine structure can be resolved by our full disk images.

times a day, to produce a sequence of images in steps of -1.0, -0.75, -0.5, 0.25, 0.0, 0.25, 0.5, 0.75, 1.0Å . These sequences provide information on any prominent radial velocity patterns. Once every 30 minutes, an image is over-exposed by an automatic exposure control so that the limb prominences are more visible. Flat field corrections are now routinely applied to Hα full-disk images, facilitating our study of filaments and prominences. Raw data from all four other stations are sent to BBSO for centralized processing.

When a filament eruption is detected, several new movies from the original 2K×2K digital data will be made to optimize the field-of-view of the targets of observation. The contrast of the filaments/prominences will be enhanced. For example, high resolution movies are made by using partial frames that cover particularly interesting filaments; and prominence movies are made by image enhancement techniques developed at BBSO. These movies will be posted on the BBSO home page promptly. Properties of each erupted filament will be studied in detail, as described below. Excellent time cadence will allow us to distinguish filament eruptions (usually associated with mass motions) from evolutionary changes that may cause a filament to disappear without erupting.

If a region is on disk center, either behind or forward spacecraft data will be used.

Most previous studies of the association of activity with CMEs were done in a traditional way – first, identify the onset of CMEs, and then locate optical and other observations to see if there is an associated Figure 3: Comparison of a CME observed on December 9, 2006 as observed by COR2/SECCHI on board

STEREO (left) and a full disk Halpha obtained at BBSO on December 11, 2006. This simulates the situation when the STEREO has a view angle of about 24 degrees away from earth’s observations. Then the center region active region’s CME will be view clearly as limb event.

solar activity (St. Cyr and Webb, 1991; Sheeley et al., 1975; Munro et al., 1979; Webb and Hundhausen, 1987). However, there are two major shortcomings in this approach: (1) CME may be associated with disk counterparts hidden behind the solar limb; (2) Some weak halo effects of disk CMEs may be overlooked, if one is not assiduously looking for them. We will study this type of association from a different perspective. We will start from Hα movies, catalog every sudden disappearance of prominence and filament, and compare it with the corresponding SECCHI movie to establish the eruption/CME association.

We will emphasize, in this proposal, a determination of the physical reason behind the association between filament eruptions and CMEs. It is essential that we obtain a large number of events to study the properties of the two kinds of filament eruptions (here we have already excluded events due to thermal heating); those associated with CMEs, and those not associated with CMEs. Based on comprehensive data comparisons, we will study the properties of each of the roughly 100 events and answer the following questions for each of them:

Does the filament have a sigmoid configuration?

Is it an AR filament or quiet sun filament?

Is it accompanied by a two ribbon flare? Based on the limited samples studied so far, a filament eruption with two ribbon flare is almost always associated with a CME.

Is its size above or a certain threshold? This threshold will be established during our statistic study. Is it a polar-crown filament? Is there new flux emergence nearby erupting filaments? Did the filament reform within certain time scale? Based on the answers to the above questions, we will find, statistically, if the properties are different for

eruptions with and without associated CMEs. If we can find important factors for which a filament eruption is associated with a CME, we may make predictions of CMEs based on the morphology of filaments that are likely to erupt.

Figure 4: Left panel: time profiles of the flare ribbon separation, filament heights measured from KSO and EIT images, and CME heights measured by Dr. Yashiro (http : //cdaw.gsfc.nasa.gov/CME_list/). The lines indicate the least-squares fits to hyperbolic functions. Right: velocity profiles of the ribbon separation, filament, and CME derived from the fits of the height profiles.

Another very important objective in our study is to follow the evolution of filament and CME height to understand the acceleration of eruption process. Again, this is improved with 3-D observations. In Figure 4, we demonstrated an example of our previous study of September 12, 2000 event, where we combine Halpha, LASCO and EIT data. For the filament eruptions, we had to assume that it moved in the direction of normal surface. Combined SECCHI and Hα observations will allow accurate determination of eruption direction and therefore accurately decide the direction of eruption.

C.2 Helicity Injection Associated with Filament Eruptions

There have been many reports on the helical structure of solar and heliospheric magnetic fields as observed in photospheric magnetic fields (Pevtsov and Canfield, 2000), coronal X-ray images (Canfield and Pevtsov, 1999), solar filaments (Chae, 2000), coronal mass ejections (Rust, 1999), and interplanetary magnetic clouds (Burlaga, 1988). So far most observational studies of the helicity of solar active regions have focused on the current helicity, defined as ∫ B · JdV, and its sign (or the linear force-free coefficient α), mainly because they can be directly inferred from photospheric vector magnetograms (e.g., Hagyard and Pevtsov, 1999). The current helicity is a measure of the topological properties, such as twist and mutual linkages of the lines of electric current (Berger and Field, 1984, Devore, 2000). On the other hand, magnetic helicity, defined as ∫ A·BdV, is a measure of twist and linkage of magnetic field lines (Devore, 2000). The magnetic helicity is a physically more useful concept than the current helicity, because magnetic helicity is fairly well-conserved in a closed volume, in the absence of boundary flows. It is considered as a robust invariant in space plasmas, such as the solar corona (Berger, 1999). However, this quantity has rarely been measured because of the difficulty in determining the topological connection of field lines in a 3D space. Instead, there have been efforts to determine its rate of temporal change in an open volume.

Since the solar corona is an open volume with the photosphere as a boundary with normal flux, the magnetic helicity can be transported across the boundary by velocity fields in the photosphere. According to Berger and Field (1984), the Poynting theorem for magnetic helicity in an open volume is given by:

\frac{dH}{dt} = 2∮ (B · A_{p})v_{z} dS − 2∮ (v · A_{p})B_{z} dS  (1)

where A_{p} is the unique vector potential of the potential field determined by the following conditions:

∇ × A_{p}· z= B_{z}, ∇ · A_{p} = 0, A_{p}· z= 0  (2)

Equation (1) tells that the magnetic helicity in an open volume can change either by the passage of field lines through the surface (first term) or by the horizontal motions of field lines (second term). The second term is particularly dominant if emerging (or submerging) fluxes are insignificant. The second term is called the shuffling term and the first is the advection term.

The BBSO group has recently studied the injection of helicity into active regions and its possible correlation with the onset of flares. Chae (2001) has developed a tool to derive the second term in the equation from the study of line of sight magnetograms. These are written in equations (3) and (4):

A_{p,x} = F T^{−1} [\frac{jk_{y}}{(k_{x}^{2}+ k_{y}^{2})} F T (B_{z})]  (3)

and

A_{p,y} = F T^{−1} [−\frac{jk_{x}}{(k_{x}^{2}+ k_{y}^{2})} F T (B_{z})]  (4)

Chae et al. (2001) studied formation of a filament, and found that the motion of footpoints provide sufficient shuffling helicity into the forming filament. Moon et al. (2002) applied the method to MDI data to study the helicity injection rate associated with several flares, and found a nice correlation between flare energy and amount of helicity injection. Figures 9 through 11 demonstrate the method of calculation of helicity injection rate for NOAA AR 8011 on January 17, 1999 (Chae, 2001.) Figure 9 is the vector map of A. Figure 10 is the horizontal flow map as derived from MDI magnetograms by local correlation tracking. Figure 11 shows the gray-scale map of −2(v · A_{p})B_{z}, which is the shuffling term of the helicity injection.

We recently studied the relationship between flares and helicity injection (Park et al., 2007). Figure 5 shows the time evolution of the integrated magnetic helicity injection together with total average magnetic flux of the four active regions. The times of X-class flares are presented by the over-plotted GOES soft X-ray light curve. These samples clearly show the helicity variations in two distinct phases: in phase I (starts half to a few days before the flare onset), the helicity accumulates at a nearly constant rate; in phase II (usually starts 1 to 17 hours before the flare onset and lasts until 1-17 hours after flare), the helicity injection rate is negligible and the magnitude of helicity remains almost constant. In all cases, we can also see a period before phase I, i.e. when the magnitude of helicity remains almost constant without helicity injection. In addition, the positive correlation between the amplitude of the flare and helicity injection rate was found.

We propose to adopt these techniques to extend the study to the evolution of helicity injection rate before, during and after filament eruptions and CMEs. We will select 10 to 20 events, which will satisfy the following criteria:

(1) Regions are close to disk center to avoid confusion arising from geometric projection.

(2) No major observed flux emergence or cancellations, so the advection term can be ignored.

(3) Filament eruptions are associated with CMEs, so that there is substantial amount of energy and mass release, as well as possible helicity transport into interplanetary space. The knowledge we gain here will also help us to understand the physical picture of helicity injection for CMEs.

Among our selected events, we will include some active region events, and some quiet region events. We understand that there are not many quiet sun events that will satisfy the above selection criteria. However, Figure 5: Time profiles of helicity injection, magnetic flux and GOES X-ray flux for 4 active regions. The

helicity is shown as crosses; the dotted lines are the GOES X-ray flux and total average magnetic flux are shown as diamonds. The flux is averaged of the absolute sum of positive and negative fluxes. we have identified three events already and will expect to identify more over the next three years. Figure 12 shows results of studies of four large flares (Moon et al., 2002). It is demonstrated that the helicity injection is enhanced around the time of the flares. Three of the events had associated filament eruptions and CMEs. This plot demonstrates the feasibility of the study proposed here.

We will use Hα full disk data to identify the time and location of eruptive events and MDI magnetograms to derive the shuffling term of helicity as described above, and demonstrated in Figures ?. Based on the studies of candidate events, we will learn if the shuffling helicity injection plays an important role in the triggering of filament eruptions. Furthermore, we will convert helicity to the injection of energy, and compare to the energy of associated CMEs. For each event, we will construct helicity injection maps as a function of time, starting two days before an events until two days after, with a cadence between half and one hour. We will be able to find if an impulsive helicity injection is important in triggering filament eruption, in the similar way as is being found for flares (Moon et al., 2002). For some selected events, we will have vector magnetograph data from BBSO which have better sensitivity, cadence and resolution than the MDI data. We will compare the calculations of helicity injection from two data sources to ensure the accuracy and reliability.

We understand that we are only considering the shuffling term in the helicity injection. During the process of selecting events for this study, we will evaluate many events to see if the magnetic flux emergence or cancelation is important in the process of filament eruption. So we will have a by-product of this study, i.e. the percentage of filament eruptions associated with flux emergence or cancellation, and their relative importance in contributing to the advection term in helicity injection. This part is difficult as we need to accurately measure vector magnetic fields as well as velocity fields. Kusano et al. (2002) developed a method of measuring helicity injection rates of both terms associated with flares. It requires the vector magnetograms. Although BBSO’s new digital magnetograph system can provide high quality vector magnetograms, an erupting filament is typically not located at the strong magnetic field part of an active region. Therefore, it is very hard to calculate the advection term for the study of filament eruptions, even with Kusano et al.’s method.

One point noteworthy is that based on the studies by Moon et al. (2002) and Chae et al. (2001), the differential rotation did not contribute a major part to the injection of helicity. The shear motion around the neutral line is the major source of helicity injection. We will find out if this conclusion is still valid for the studies of helicity injection associated with filament eruptions.

C.3 The Relationship between the Rate of Magnetic Reconnection and Solar Eruptions

Magnetic reconnection is considered to play an important role in driving the solar eruptions (filament eruptions and CMEs) as well as giving rise to intense heating that accounts for the associated flares (Lin 2001). Numerous theoretical models for solar eruptions and the evolution of the reconnecting current sheet (RCS) have been developed (e.g., Van Tend & Kuperus 1978; Kaastra 1985; Martens & Kuin 1989; Forbes & Priest 1995; Amari et al. 1996; Lin & Forbes 2000; and see Lin 2003 for review). The rate at which the magnetic flux is converted into the diffusion region, in terms of the electric field E_{rec} in the RCS, can be inferred by measuring the flare ribbon expansion speed V_{r} and the normal component of the magnetic field B_{n} swept by the flare ribbons: E_{rec} = V_{r}B_{n} (Forbes & Priest 1984). More generally, Lin and Forbes (2000) considered the rate of photospheric magnetic flux change φ_{rec} instead of the electric field in the region of newly closed field lines. The flux change rate can be evaluated by the following equation (Lin & Forbes 2000):

φ_{rec} = ∫ E_{rec}dl = \frac{∂}{∂t}∫ B_{n}da ,  (5)

where dl is the length along the RCS and da is the newly brightened area swept by the flare ribbons.

Figure 6: Top panels: snapshot of an M1.0 flare in Hα observed by KSO on 2000 September 12 and corresponding MDI magnetogram. Plus signs are trajectories of the ribbon fronts; colors from purple to white indicates a 2 hr time lapse from 11:00 to 13:00 UT. Bottom panels: Left: magnetic reconnection rate in terms of E_{c} (red) and φ_{c} (blue) in comparison with the time derivative of the GOES soft X-ray emission (dark) for the M1.0 event on 2000 September 12. Right: magnetic reconnection rate in terms of E_{c} (red) and φ_{c} (blue) in comparison with the microwave light curve at 6.6 GHz (dark) for the X1.6 event on 2001 October 19 (Qiu et al., 2004).

E_{rec} and φ_{rec} provide a measure of the reconnection rate inside the current sheet, and the measurement itself is not dependent on specific models. In this manner, we are particularly interested in the observational signatures of two-ribbon flares that are associated with solar eruptions. We can infer the reconnection rate and compare it with the rising motion of flux-ropes (such as erupting filaments or CMEs). Essentially, the reconnection rate and the flux-rope acceleration focus on important features of the flare. Their connection, therefore, provide a stringent test of the scenarios in which the ejection of solar material and magnetic reconnection are physically linked.

As an example, Figure 6 shows snapshot of an M1.0 flare observed at by Hα Network on 2000 September 12 and corresponding MDI magnetogram, with evolution of flare ribbons marked. It also shows temporal correlation between the flare emissions and reconnection rate for the event on 2000 September 12 and the X1.6 flare on 2001 October 19.

We have been using GHN Hα full-disk images as the primary data source to trace the flare ribbon expansion and the filament rising motion. These Hα observations are suitable because of their high-cadence (one or more image frames per minute), high-resolution (approximately 1 arcsec/pixel) and superior detail and contrast (Martin 1989). The CME height-time data is provided by the LASCO web site that has been compiled by the GSFC group (http : //cdaw.gsfc.nasa.gov/CME_{list}/). The velocity and acceleration of the Figure 7: Top-left: Scatter plot of maximum filament acceleration vs. maximum φ_{rec}; Bottom-left: Scatter

plot of maximum filament acceleration vs. maximum E_{rec}. Top-right: Scatter plot of CME velocity at C2 vs. maximum φ_{rec}; Bottom-right: Scatter plot of CME velocity vs. maximum E_{rec}. Error bars attached to each sign indicate the uncertainty of the measurement (Jing et al., 2005).

filaments and CMEs can be derived numerically as the first and second derivative of corresponding height with respect to time. Initially, MDI magnetograms will be used to measure the longitudinal component of the magnetic fields, which approximate the normal component of the magnetic field B_{n} because only those flares that occur near the disk center will be under our study. BBSO IRIM data will provide more accurate calibration and measurement of true field strength.

From a sample of 13 well observed two-ribbon flares, that are associated with filament eruptions or CMEs, acceleration of erupting filaments is mainly in the range of 0.05 − 0.4kms^{−2} , up to 3kms^{−2} . The maximum E_{rec} and φ_{rec} mostly occur in the range of 0.2 − 5V cm^{−1} and 0.5 − 6 × 10^{18}Mxs^{−1} , respectively. In the case of the extremely dramatic flare on October 29, 2003, E_{rec} reaches a magnitude of about 38V cm^{−1} . We found the electric field is generally comparable with most observational results (Poletto & Kopp 1986; Wang et al. 2003, 2004; Qiu et al., 2004) and simulation results (Martens & Kuin 1989; Lin 2002; Cheng et al. 2003). The electric field strength found above implies strong heating and particle acceleration which are responsible for the high energy emissions (Cheng et al. 2003; Qiu et al. 2002, 2004). The left two panels of Figure 5 displays scatter plots of the reconnection rate (namely, E_{rec} and φ_{rec}) versus the acceleration of filaments in a logarithmic scale. All values refer to their maximal that were found on each events. Certain linear correlation between filament acceleration and reconnection rates is demonstrated, but need more events to confirm.

Due to the temporal-spatial gap between the disk observations of the flares that involve magnetic reconnection and LASCO observations of CME fronts, it is less meaningful to compare the magnetic reconnection rate and the CME acceleration. Instead, we use the velocity of CMEs at C2 to examine the Figure 8: Scatter plot of CME velocities vs. to-

tal reconnection flux for 13 events. Dark and gray colors indicate events associated with erupting and non-erupting filaments, respectively. The dashed guideline shows the least-squares linear fit to the data points (Qiu and Yurchyshyn, 2005).

relationship between CMEs and flares in that the velocity is the cumulation of the acceleration at its early stage (Vel_{CME}= ∫ Accel_{CME}dt) and, therefore, in some way conveys information about the acceleration during that stage. A graphical display of the CME velocity versus the reconnection rate is presented by two panels on the right side of Figure 5. The figure suggests some weak correlation between them.

More recently, Qiu and Yurchyshyn (2005) explored relationship between the total reconnection flux estimated from flare observations and the velocity CMEs. It is shown that the velocity of CME is proportional to the total reconnection flux, with a linear cross-correlation of 89% and confidence level greater than 99.5%. This result confirms the importance of magnetic flux transferred by magnetic reconnection in the early stage of fast CMEs. Figure 6 shows their results.

In summary, with the limited sample, we had found certain magnitude scaling relationship between the magnetic reconnection rate and the kinetics of erupting filaments and CMEs. These pilot results need to be certified by including more events to improve the statistical significance. We propose to continue this study by including many more events. In order to achieve this goal, we will improve our automatic program to align Hα images and MDI magnetograms, count all brightened pixels swept by flare ribbons in a sequence of consecutive Hα images and then map them to the co-registered longitudinal MDI magnetograms. This program, mainly based on “region growing", “adaptive boundary-base" and “morphological closing" methods, makes batch processing possible. We anticipate to study over 300 events to firmly understand the link between the magnetic reconnection and kinetics of solar eruptions. Moreover, it is important to determine whether the relationship is independent of the magnetic configurations of source regions and the structure of CMEs. If and when the correlation relationship can be solidly established, we will provide a real-time prediction of filament accelerations and its associated CME speeds, based on the real-time monitoring of flare ribbon separation. These will be more important for these events that filament and CME kinetics can not be derived easily from direct measurements (e.g. certain fast filament eruptions and Halo CMEs).

The tasks for this part are summarized as following:

1. Implement real time tracking of two ribbon flare evolution.

2. Statistical study of relationship between magnetic reconnection and filament/CME kinetics.

3. Implement the real-time monitoring and forecasting web tool.

D Personnel and Budget

This research project will be carried out primarily by the e-member research team: PI Wang, Co-Is Drs. Yurchyshyn and Jing. Dr. Yurchyshyn is a research professor at BBSO Dr. JING is a post-doctoral research fellow at NJIT. Although the entire team will work closely, each individual will focus on particular area. Dr. Jing will be primarily responsible for calculating magnetic reconnection rate, as she has substantial experience in this area. Dr. Yurchyshyn will primarily carry out statistical studies of filament eruptions and CME onset relations. Wang will lead the study of helicity injection with an NJIT paid graduate student.

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