5 Instabilities and Prominence Disappearance

Prominences end their lives either through a slow process of decaying emission, or through a more sudden disappearance (disparitions brusques) triggered by a thermal or dynamical instability. The decay of emission can also be transient: the disappearance of a filament in a monochromatic image sampling the cool plasma (Hα for example) can be the result of the change of its temperature as a consequence of a transient heating injection. Thus, the filament will become visible in the hotter EUV lines, but eventually will go back to its original thermal condition. A more violent phenomenon is the prominence eruption, which implies an explosive rearrangement of the magnetic structure and its ejection into the extended corona (see Section 5.4). Nevertheless, intermediate dynamic phenomena are also observed, such as partial or failed eruptions. Prominence eruptions are associated with flares (impulsive release of energy, Benz, 2008) and coronal mass ejections (see Figures 31*, 32* and Schwenn, 2006), indicating that a common magnetic structure (the filament channel) links these energetic phenomena.

5.1 The sources of instabilities

Prominence destabilization has been observed in connection with several phenomena, including magnetic flux emergence, local and large scale photospheric motions with their transport of magnetic flux, and a remote flare that initiates wave disturbances (e.g., Wang et al., 2001).

The photospheric motions can be responsible for the shearing of magnetic field and flux cancellation through collisions of opposite polarities. Both photospheric motions and flux emergence perturb the existing magnetic field whose changes are considered to be the energy source for the eruptions. Magnetic reconnection is one of the primary ways this energy is released.

Flux emergence, small-scale surface flows, and reconnection associated with filament eruption can easily be identified in the data. Various cases will be presented in Section 5.4 in the context of CMEs. The relation between filaments and large-scale photospheric motion, on the other hand, is well studied by modeling and numerical simulations (Mackay et al., 2010), but is rarely found in observations.

However, a nice observational example of this scenario is shown in Figure 26* from Roudier et al. (2008*). These authors followed the time changes of the topology of the flow around and under a filament before and after its eruption. They found important changes in the horizontal component (which is associated with the differential rotation) in the region where the filament eruption started. Such a change was also associated with the presence of north-south streams and to a change of the orientation (toward the streams’ direction) of this part of the filament. From these observations the authors concluded that the flow dragged the filament footpoints, stretching the coronal magnetic field and enabling the subsequent instability. Further investigation is needed to confirm this picture for more observed events, and to include similar processes in existing CME models (e.g., Aulanier et al., 2010).

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Figure 26: Map of the amplitude of the change of the velocity vector direction. The closer this quantity is to 1, the better the alignment between two vector fields. Black areas indicate the location where a lot of changes between horizontal velocities take place. The filament observed before the eruption is superimposed. The hexagon indicates the location where the filament eruption started. Image reproduced by permission from Roudier et al. (2008), copyright by ESO.

5.2 Observed precursors of eruptions

The disappearance or eruption of a prominence is generally preceded by an activation phase. This shows up as an increase of mass or turbulent motion inside and at the structure footpoints (for example, as shown in Figure 27*), fragmentary brightening and fading due to sporadic heating events (e.g., Ofman et al., 1998; Schmieder et al., 2000; Kucera and Landi, 2008), EUV brightenings due to reconnections, deformation and helical rotation of the filament (Sterling et al., 2011*), and the start of lifting off (e.g., Gilbert et al., 2000*; Engvold et al., 2001; Sterling and Moore, 2004*; Bemporad, 2009*; Liewer et al., 2009*). Figure 28* shows a further example of these pre-erupting signatures: the Hα intensity variation in a segment of a filament, which will start disappearing at around 12:30 UT.

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Figure 27: Hα intensity and velocity images observed using MSDP instrument at Meudon Solar Tower on 21 May 2008, at 09:00 UT. The left panel shows the line center intensity while the right panel shows, in grey scale, the dopplergram overlaid by contours of the filament. The white arrows mark the maximum (white) and minimum (black) velocities, respectively, +800 m/s and –200 m/s. The white circle indicates upward and downward flows at the end of one foot of the filament. Elongated areas of blue and red-shifts (marked by the arrows) suggest some twist along the filament body. Image reproduced by permission from Gosain et al. (2009*), copyright by Springer.
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Figure 28: Light curve and stackplots in X and Y directions showing the variation in the Hα filament intensity on 19 May 2007 prior to the filament eruption at around 12:50 UT. The top panel shows the total emission, the middle is the variation in Y intensity and the bottom is the variation in X intensity. Image reproduced by permission from Bone et al. (2009), copyright by Springer.

Prominence oscillations are also observed in the pre-erupting phase of a filament (e.g., Isobe and Tripathi, 2006; Pintér et al., 2008). An example is given in the movie of Figure 29* obtained from the STEREO/EUVI data. Contrary to the quiescent case, where only a few periods of oscillations are observed, the number of periods in an activated filament may be numerous, as observed by Chen et al. (2008*) using SOHO/SUMER data. These long-lasting oscillations were interpreted as the signature of the repetitive action of the exciting force. The authors also noted an association between the oscillation and the presence of surges and siphon flow in the Hα images. Such events were interpreted as signatures of reconnection of emerging flux with the existing coronal field (Figure 30*), as we will discuss later in this section.

Martin (1980), Gopalswamy et al. (2006), and Schrijver et al. (2008*) review in further detail the pre-erupting conditions and ejection phase.

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Figure 29: Liftoff and eruption of a filament observed with STEREO/EUVI. An oscillatory motion is clearly visible during the rising phase before the eruption. Courtesy of STEREO/EUVI consortium.
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Figure 30: Left panel: evolution of the C i 1118.45 Å intensity along the SUMER slit; Middle panel: same for S iii/Si iii 1113 Å; Right panel: evolution of the Dopplergram along the SUMER slit observed at S iii/Si iii 1113 Å. Image reproduced by permission from Chen et al. (2008), copyright by ESO.

5.3 Eruption phase

In most cases (particularly for quiet region structures) the initial rise phase of the filament is quite slow (1 – 15 km s–1) and can last a few hours, while this phase can be much shorter (10 min) for active region structures (Sterling and Moore, 2004*; Williams et al., 2005*). This is followed by an acceleration phase in which the filament may reach about 100 – 1000 km s–1, together with the rest of the expelled coronal material to form the CMEs (e.g., Schrijver et al., 2008; Gosain et al., 2009*). In a final stage the structure expands at a nearly constant speed or may decelerate. Depending on the magnetic environment, the prominence eruption can be either almost symmetric or asymmetric with respect to its endpoints (Liu et al., 2009a).

Partial or failed eruptions of prominences have an initial phase similar to that of a full prominence eruption up to the acceleration phase, which lasts for a shorter time. The prominence is then observed to decelerate, reaching a maximum height, while prominence material simultaneouly drains back to the solar surface. Failed eruptions sometimes are associated with CMEs. In this case the remaining filament reforms soon, to finally be completely ejected in the CME (Zhou et al., 2006). Other times there is no sign of a CME or opening of the coronal magnetic field. Often these events are associated with compact flares. Similar to full prominence eruption, failed eruptions show properties that can be reproduced by models that use different destabilizing mechanisms, as we will see in Section 5.4.

In the off-limb images showing the eruption phase, the prominence generally can be better isolated against the darker coronal background than on the disk. Often in this case the prominence structure turns out to be composed of bright (emitting) and dark (absorbing) parts highlighting a twisting topology, suggesting the presence of a magnetic flux rope (see Figure 31* and, e.g., Patsourakos and Vial, 2002). As discussed earlier, it is not yet clear whether this topology results from the instability and following eruption or existed beforehand.

5.4 Eruption of prominences associated to flares and CMEs

Since the SkyLab era, several statistical studies have addressed the relation between filament eruptions, flares, and the ejection of CMEs (e.g., Munro et al., 1979; Webb and Hundhausen, 1987; Gilbert et al., 2000*; Zhang et al., 2001); the results range from a 10% (Hori and Culhane, 2002) to a 90% (Gopalswamy et al., 2003) association. More recently, Al-Omari et al. (2010) used the NGDC filament catalogue18 and the SOHO/LASCO CME catalogue to feed the SVM learning algorithms and study the association between CMEs and filament eruptions in the period 1996 – 2001. They found that about 4% of the observed CMEs were associated with filaments. Aschwanden et al. (2009) made a recent survey of flares using the first two years of STEREO/EUVI data (November 2006 – November 2008) and claim that only 3% of them were associated with filament eruptions. Taliashvili et al. (2009) investigated 42 sudden disappearances of quiescent prominences during solar minima and classified them as thermal or dynamic. Their main results show that 50% of the thermal and 60% of the dynamic disappearances were associated with coronal hole proximity and/or CME occurrence.

These diverse results obtained before the launch of SDO indicate difficulties in defining the strength of the association among filament eruptions, CMEs, and flares. Different results may originate from different selection criteria adopted for the events, differences in the solar height of the FOV of the observations, differences in the chosen temperature coverage, and the presence or absence of simultaneous observations of CMEs, flares and filaments. More systematic selection criteria need to be assured in the future, together with full temporal and temperature coverage of the events (even though SDO data are now improving much on these latest two aspects).

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Figure 31: A SOHO/Lasco C2 image of a CME. The central bright helical structure is identified with the erupting filament. Credits: SOHO (ESA & NASA).

When a filament eruption is associated with a CME, the filament is identified as the brightest part of the CME in white light images (for example, see Figure 31*). Observationally, a CME structure is often very similar to the large scale structure in which we find the filament: a bright leading edge, a dark cavity, and a bright core (the erupted filament) (Hundhausen, 1999). This suggests that in many cases a large fraction of the whole filament-channel magnetic structure is destabilized and erupts (see Figure 32*).

To assess how much of the coronal environment is involved in the prominence eruption, how the destabilization is triggered, and how it evolves, we need good time sampling of the event at multiple wavebands. As a result of recent advances, some common observational signatures of the pre and post-eruption have now been identified, as sketched in Figure 32* and described below. In addition, other signatures may be present and may change from case to case (these will also be mentioned later). The diagnostics and timing during the eruption of all these properties have been used to constrain and develop different models, which are summarized next.

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Figure 32: Left panel: a Solar Maximum Mission archive image showing the principal features of a CME Image reproduced by permission from van Driel-Gesztelyi and Culhane (2009), copyright by Springer. Right panel: schematic view of the CME features. Image reproduced by permission from (Forbes, 2000*), copyright by AGU.

5.4.1 Models for eruptions

Different eruption models predict that different triggering mechanisms are responsible for the energy release through the eruption of the CME, after sufficient magnetic energy has been accumulated up to a critical point. The energy release can be the result of instabilities invoking ideal or resistive processes. Models associated with the first group, for example, invoke helical kink or torus instabilities of a flux rope (e.g., Démoulin and Priest, 1989; Amari et al., 2003a; Fan and Gibson, 2003; Török and Kliem, 2005*; Williams et al., 2005*). Kink instability occurs when the twist reaches a critical value, while for torus instability the background magnetic field gradient plays a key role. In both cases, the flux rope then deforms and rises up to its ejection.

Those models that evoke instabilities by flux cancellation and reconnection belong to the second group (e.g., van Ballegooijen and Martens, 1989). In “tether-cutting” models (Moore and Roumeliotis, 1992; Moore et al., 2001), the instability and reconnection can originate first below the filament in a bipolar sheared system. Here the filament and the formed flux rope are pulled outwards by the release of magnetic tension associated with the underlying reconnection, simultaneously powering a flare below the reconnection site. In the “breakout” model (Antiochos, 1998; Lynch et al., 2008; Karpen et al., 2012*), which requires a multipolar magnetic flux system (e.g., a delta spot) and a sheared filament channel, two reconnections are involved: “breakout reconnection”, which removes the overlying field restraining the ejection, and “flare reconnection” beneath the filament, which forms the CME flux rope as in the tether-cutting paradigm.

Within these models, explanations also have been developed for the partial or failed eruptions of filaments (e.g., Gilbert et al., 2000; Gibson and Fan, 2006*; DeVore and Antiochos, 2008, and references therein).

Most observed events can be interpreted with one or more of these models. For a more extensive overview of CMEs and prominence eruption modeling see, e.g., Forbes (2000), Forbes et al. (2006) and Chen (2011).

5.4.2 General observational properties

The general observational aspects that can be seen during an eruption, which should be reproduced by models, are sketched in Figure 32* and reviewed by Webb and Howard (2012*). Before and during an eruption, EUV and HXR local brightenings as well as radio burst emission, are present and are considered signatures of magnetic reconnection.

The eruption models agree that a formed flux rope erupts once it exceeds the instability threshold, which forces the overlying magnetic field to open. However, they differ in the location where the different events happens, as well as their temporal sequence (the formation of a current sheet, for instance). Magnetic equilibrium is eventually recovered as a consequence of reconnection in the current sheet. Below the erupting large-scale structure, the core flaring region is filled with a post-eruption arcade whose loops are observed in the EUV and SXR images. The footpoints of these loops are identified in two bright ribbons seen in Hα and UV-EUV images and in HXR sources. The ribbons separate in time and their inner edges mark a transient coronal hole. This separation is interpreted as a consequence of the rising of the magnetic reconnection point at the current sheet formed below a CME. HXR sources observed between the top of the loop arcade and the bottom of the erupting flux rope (see Figures 32* and 33*) have been identified as either the reconnection site or the termination shock of the downward reconnection jet. While the location of the reconnection-site HXR source is expected to move up as the separation of the ribbons from the PIL increases (implying a rising reconnection point), in some cases it is observed to move along it (Grigis and Benz, 2005; Tripathi et al., 2006*). This aspect has been interpreted as the consequence of an asymmetric eruption of a filament, i.e., when one end of the filament erupts before the other. However, the signature of these HXR sources is still debated. Radio emission as bursts, particularly type II, III and IV are also detected during CMEs. Type II bursts are tracing shock waves, type III are due to fast electrons escaping the corona along the magnetic field, and type IV are due to non-thermal emission, typical of the regions where magnetic reconnection is at work. Type III bursts are sometimes detected a few hours before onset of the eruption (see Webb and Howard (2012) for details).

In addition to the timing of the coronal brightenings described here (examples relevant to model predictions are given below), another element that can provide insight into the eruption initial phase is studying and comparing the kinematics of different components of the global eruption, in particular the height and acceleration of the moving filament and the CME leading edge. For the models that predict a whole flux-rope instability, the eruption of the rope (the magnetic support embedding the prominence plasma) is almost uniform, leading to synchronized motion of both prominence and CME leading edge (the upper part of the flux rope). This is the case, for example of Török and Kliem (2005*).

For those models that invoke resistive instabilities, such as the breakout and tether-cutting models, we expect less synchronization in the initial phase of the whole CME eruption. In a detailed study of a simulated breakout CME, Karpen et al. (2012) found that the prominence-carrying portion of the structure undergoes a short stage of Alfvénic speeds during the classic impulsive phase of the CME/eruptive flare, but then travels more slowly than the CME front.

Results of observational studies focusing on relative kinematics have shown that there is not a unique picture. Most of the 18 cases studied by Maričić et al. (2009), for example, revealed that the motions of the leading edge of the CME and the erupting prominence were very similar and almost co-temporal. However, they also found exceptions where the prominence acceleration started earlier (< 40 mins) than that of the CME leading edge, and others where the opposite happened (with a similar delay) prior to eruption.

A distinction between these two latter cases is difficult, however, so that this kind of analysis should be accompanied by a more reliable study on the whole erupting environment. Among other factors, the relative delay of the two fronts can depend on the location within the magnetic configuration where the instability starts (e.g., the position of the reconnection site). To clarify this situation, the full-time history of the reconnection sites should be identified and recorded. In addition, the possible natural expansion of the erupting flux rope, which affects the synchronization of the two fronts, should be taken into account and, in some way, disentangled from a lack of synchronization due to the initial conditions of the eruption.

Furthermore, we must consider the difficulty of these kinematic studies arising from the fact that the CME emission is quite faint compared to the background corona. Multi-instrument observations (such as those from the STEREO mission) can follow the initial phase of the events using the overlapping field of views of the disk imagers (showing the solar disk and only a small fraction of the off-limb corona in the EUV) and coronagraphs (showing only the off-limb corona in the optical band). However, up to now this overlapping region is quite small and the emissions recorded by these instruments are due to different physical processes. Such difference should be taken into account in the data interpretation. However, additional information is obtained through the technology of the new generation of ground-based coronagraphs (e.g., COSMO K-cor at NHO19) and space instruments (e.g., the SOLAR Orbiter/METIS coronagraph and EUI imager), where the overlapping field of view is increased.

Finally, in terms of localization of the initial instability, the additional data provided by the latest solar mission, SDO, with AIA’s high temporal (10 s) cadence and multi-temperatures images, together with the magnetogram data provided by HMI, are now helping to solve this problem.

At present, we have collected enough examples of prominence eruptions and CMEs to realize that it is very difficult to associate all of them with well-established (tested) models. The corona is highly structured and dynamic over large spatial scales, so that unusual events may always occur, or external factors may have a role in the dynamic evolution of the eruption. For example, a series of observations has shown that there may be an offset between the direction of the filament eruption and the CME expansion, an indication that a complex magnetic configuration was present before the instability (Gopalswamy et al., 2000) and/or the pre-existing large-scale coronal structure influenced the CME expansion. Recent results indeed indicate that the pre-existing coronal streamer above the filament may constrain the filament and CME expansions after their eruption (Bemporad et al., 2005; Jiang et al., 2009). An interesting point has recently been raised by Liu (2008) and Liu et al. (2009c) on the role of the strength and asymmetric background magnetic coronal field in constraining the filament eruption. Using simple potential field modeling and observational cases, they showed that an asymmetric field can produce a stronger confinement in the trajectory of a filament eruption than in an symmetric case, also causing failed eruptions. These observed events suggest that models predicting the CME expansion need to evolve to include a more realistic coronal environment.

Magnetic field reconnection.
We summarize some recent research that identifies the EUV and X-ray brightenings with the signatures of reconnection associated with eruptions. These cases are meant to illustrate the different interpretations that may be given to various observed events, showing that it is not always clear whether reconnections are at the origin of the eruption or its consequence (for example, see Forbes and Lin, 2000).

One of the key approaches to answering this question depends on our ability to identify the location and timing of the EUV and X-ray brightenings. Their position within the magnetic structure, inferred from the photospheric magnetic field and its extrapolation (when available), can help distinguish among the various magnetic models for the eruption. For example, the presence of a photospheric multipolar region and pre-eruption reconnection far above the filament channel are consistent with the breakout model, while a simple photospheric dipolar region is consistent with a tether-cutting configuration. Both models predict signatures of reconnection below the filament during the impulsive phase and eruption, with the presence of sigmoids (an ‘S’ shaped SXR structure).

An example showing the difficulty in data interpretation for the brightenings is given by Sterling et al. (2001b*). They observed the gradual lift off of a quiescent filament, followed by a more rapid phase ending with its eruption in correlation with a flare. The absence of SXR emission below the filament prior to its eruption suggested to the authors a gradual initial phase consistent with breakout reconnection. However, the lack of high cadence data that could catch the brightenings/dimmings associated with early reconnection implied that they could not rule out the tether-cutting scenario (see also Sterling and Moore, 2004). This is a clear case showing the importance of full coverage and high cadence data in order to fully capture the ongoing dynamics. Similar cases to those observed by Sterling et al. (2001b) were also found in other observations by the same authors (Sterling and Moore, 2001; Sterling et al., 2001a); generally they support the presence of a mechanism that induces filament destabilization prior to the explosive phase.

EUV and X-ray brightenings have also been interpreted as signatures of reconnections due to, for instance, filament kink instability. This is the case of Liu and Alexander (2009*) who detected the formation of multiple hard X-ray sources visible at the interacting filament legs in the first phase of prominence activation, and beneath the filament arch during its rising phase (Figure 33*). Liu et al. (2009b) instead, found their data consistent with both tether-cutting and kink instability as the triggering mechanism for the observed flare and filament eruption. They observed thermal hard X-ray emission close to one filament leg and the subsequent formation or increase of filament twisting.

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Figure 33: Prominence eruption interpreted as the result of a kink instability. TRACE 195 images overlaid onto RHESSI X-ray contours at 10%, 20%, 50%, and 80% of the peak intensity at 1220 keV (gray), and 2040 keV (black). Coronal sources are marked by white arrows. The integration time for each RHESSI image is 12 s around the time of the corresponding TRACE image. Image reproduced by permission from Liu and Alexander (2009), copyright by AAS.

Other observations detect reconnection sites both above and below the filament during different phases of the eruption. This was, for example, the case of the failed active-region filament eruption followed by Alexander et al. (2006) with TRACE and RHESSI, that was associated with an M2 flare but not with a CME. The initial rising phase of the event was interpreted as a flux rope eruption due to a kink instability (Török and Kliem, 2005). By studying the changing positions of the HXR sources with respect to the evolution of the lifting filament, the authors inferred that a weakening, through reconnection, of the coronal magnetic field occurred above the filament, while interaction/reconnection of the two footpoints was going on at the coronal base.

In the case of Nagashima et al. (2007), on the other hand, a sequence of small flares close to the filament footpoints was identified at the time the filament was slowly lifting off (40 h before the eruption). This was followed by the filament eruption and a CME. This sequence of events was interpreted in terms of reconnections in the photosphere that produced changes in the magnetic configuration above the filament of the overlying loop arcade, leading to the filament eruption.

As we have seen, the location of the brightenings can be different for different erupting events. A systematic study to interpret this aspect was done by Tripathi et al. (2006*) using data from seventeen events, relating the separation of the two-ribbon flare with the propagation along the neutral line of the SOHO/EIT 195 channel brightenings associated with filament lift off. They could distinguish symmetric from asymmetric eruptions. In the first case the brightening started from the center of the filament and propagated toward its end, while in the latter the brightening moves from one end to the other of the erupting filament. A sketch of their 3D picture of the two eruptions is shown in Figure 34*.

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Figure 34: Schematic diagrams showing asymmetric (left panel) and symmetric (right panel) eruption of a filament. The arrows indicate the direction of propagation along and separation away from the polarity inversion line of the SOHO/EIT 195 channel brightenings. Image reproduced by permission from Tripathi et al. (2006), copyright by ESO.

In conclusion we can say that a large amount of data has been published in which EUV and X-ray brightenings are associated with reconnection and eruptions. They are interpreted as the signature of conversion of magnetic energy into kinetic energy and/or heat. They are also generally observed during coronal magnetic-field reconfigurations, such as twisting of the filament structure and changes of the large-scale magnetic field. The interpretation of such brightenings in the frame of models for energy ejection is not unique. The different existing configurations of the magnetic field, the long-term energy buildup by photospheric motions and by the emergence of a new magnetic field, and other factors mentioned above all contribute to the wide variation in observable signatures.

Magnetic flux emergence and cancellation.
As mentioned, simulations have shown that differential rotation coupled with surface flows contribute to the transport of photospheric magnetic-field flux and may cause field-line shearing, reconnection and/or magnetic flux rope formation (for the theory of flux rope and helicity see, for example, Berger, 1984; House and Berger, 1987; Rust, 2003). We give here a few examples that show that the reconnection events can result from flux emergence or cancellation associated with diffusion through the photospheric surface (Feynman and Martin, 1995; Muglach et al., 2009). A recent statistical work by Yan et al. (2011) shows that flux emergence and cancellation have about the same effect on active-region filament destabilization.

Chen and Shibata (2000) proposed a model where destabilization and reconnection within the filament channel or at the outer edge of the channel were initiated by nearby flux emergence. Data analyzed by Sterling and Moore (2005) were consistent with predictions of this model during both the initial gradual liftoff of the filament and the eruption.

Even though several models can predict partial-filament eruptions, the event studied by Tripathi et al. (2009a) had several elements consistent with the partially-expelled flux rope (PEFR) model by Gibson and Fan (2006), in which a flux rope emerges and becomes unstable, reconnecting with the surrounding existing field. They observed the change of the magnetic structure from a sigmoid into a new sigmoid surrounded by a cup shaped arcade, and the formation of a dimming area outside the eruption source region.

Mackay and van Ballegooijen (2006) show that flux cancellation at the PIL between bipoles may result in the formation of a flux rope, its elevation and consequent reconnection below it, and subsequent ejection. Amari et al. (2010), for example, also showed that such cancellation may transform a sheared arcade into a stable or unstable flux rope. These ideas have been supported by data on several eruptive events (e.g., Green and Kliem, 2009; Green et al., 2011; Sterling et al., 2010) and are beginning to be confirmed by the SDO/AIA data. In addition, Sterling et al. (2011) observed a case of flux cancellation below the filament with successive building up of the flux rope. The eruption was preceded by UV brightenings (attributed to reconnection) consistent with those described by the tether-cutting model.

Some of these models rely on the efficiency of photospheric diffusive motions to drive reconnection, but observational evidence for this process is sparse (for example, see the case shown in Section 5.1) and therefore still inconclusive. The event studied by Schmieder et al. (2008), for instance, show no evidence of reconnection or flux emergence (even though the authors cannot rule out their presence), while a decrease of the magnetic field strength in the network surrounding the structure was inferred during the rise of the filament. The authors suggested an interpretation in line with the presence of the photospheric diffusion process.

Several studies also suggest that the transportation of minor polarities (where the filaments barbs root) toward the PIL can lead to their cancellation and to magnetic destabilization (Lin et al., 2005; Gosain et al., 2009*).

5.4.3 Magnetic helicity

Similar to other coronal structures, the magnetic envelope of filaments (e.g., the flux rope or sheared arcade) may present a helicity that evolves in time (see also Section 2.4.1). In particular, flux emergence and photospheric flows can change the magnetic helicity in a filament channel, trigger its instabilities, and set off the ensuing filament eruption.

An increase in helicity due to flux emergence was identified by Williams et al. (2005), who suggested that this phenomenon triggered a kink instability and the following filament eruption. Similarly, Romano et al. (2009) associated the increase of helicity in an observed filament with flux emergence. In particular they identified the filament footpoints as the location of helicity injection into the corona, which led the structure to take on a sigmoid shape and then to erupt.

The role of photospheric flow in the transport of helicity from the photosphere to the corona was investigated by Romano et al. (2005*). This motion was associated with an M6.3 flare and filament eruption on 15 June 2001 in the active region NOAA 9502. They showed that, during the ‘S’ shaped filament activation phase, a horizontal counterclockwise motion at both filament footpoints was associated with an increase of magnetic helicity change rate (negative, in their case as shown in Figure 35*), without any significant magnetic flux variation (contrary to their later work mentioned in the previous paragraph). In addition, an impulsive variation of the helicity change rate was observed at the beginning of the filament eruption. This was interpreted as a signature of flux cancellation or expansion in the corona caused by the flare. In this work the important role of the horizontal motion in the helicity transport is obvious.

Flux emergence and photospheric motions are other factors that possibly trigger instabilities, as they may lead to a change of sign of the pre-existing helicity.

Evidence for these processes was presented by Su et al. (2005), who studied a ‘U’ shaped Hα filament several hours prior to its eruption. During this period they noticed an inversion in the chirality of one barb, which could have been associated with a local change of the helicity of the filament. They suggest this change could be the origin of the filament destabilization, but they could not identify the origin of the chirality change.

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Figure 35: Accumulated change of the magnetic helicity as a function of time in AR 9502 on June 2001, prior to a flare and filament eruption. t = 0 corresponds to 13:00 UT on June 14. The error bar represents the standard deviation of the signal. The vertical lines indicate the start and end of the flare. Image reproduced by permission from Romano et al. (2005), copyright by ESO.

Another phenomenon often visible during the lift off of a filament is its rotation in a direction that appears to be correlated with the chirality of the magnetic environment. In particular, Green et al. (2007*) and Rust and LaBonte (2005) found several cases that were consistent with the picture of a filament contained in a magnetic flux rope that twists as it erupts. The twist was in the same direction as the internal twist of the flux rope, indicating the conservation of the helicity.

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Figure 36: Big Bear Solar Observatory Hα images of active region 9077 on 19 July 2000 showing the counterclockwise rotation of the filament. The TRACE 1600 Å image shows the formation of flare ribbons which form a reverse-S shape and indicate negative helicity. Image reproduced by permission from Green et al. (2007*), copyright by Springer.

5.5 3D reconstruction of erupting filaments

To conclude this review, it is worth mentioning the new insights on prominence eruptions and CMEs arising from the use of the 3D reconstruction of STEREO/SECCHI images (Howard et al., 2008). The STEREO mission is composed of two identical satellites. The payload includes SECCHI, a suite composed of the Extreme UltraViolet Imager (EUVI), inner (COR1) and outer (COR2) coronagraphs, and inner (HI1) and outer (HI2) heliospheric imagers. Since 2006 the two satellites have been orbiting the Sun, one ahead and one behind the Earth, allowing a stereoscopic vision of the Sun (see the cartoon in Figure 37*). Here a few examples are given to show that new constraints can be obtained on the propagation direction, timing between the filament acceleration vs. CME acceleration, 3D morphology etc. Stereoscopic and tomographic methods are reviewed by Aschwanden (2011), Bemporad (2011) reviews in detail 3D prominence reconstruction techniques and presents several results, while Mierla et al. (2010) reviews 3D CME reconstruction. This new approach to the study of filaments and CME eruption will certainly provide more insights in the future.

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Figure 37: Cartoon showing the geometry of simultaneous observation from STEREO A and B of a prominence eruption. Given the two different projected altitudes (RA and RB) and latitudes (ϕA and ϕB) of the same point P, knowing the angular distance γ between the two spacecrafts, it is possible to derive via triangulation the 3D coordinates of this point. Image reproduced by permission from Bemporad (2009*), copyright by AAS.

Among the first studies on this subject, Liewer et al. (2009) and Gissot et al. (2008) independently reconstructed the 3D configuration of the asymmetric eruption analyzed by Li et al. (2008). The two studies used different stereoscopic techniques (stereoscopic tiepointing and triangulation for the former, and the Optical-Flow Algorithm for the latter), and their results were in good agreement. From these reconstructions they could estimate the direction and speed (≈ 100 km s–1) of the prominence, and establish that it was below that derived for the associated CME.

Using a 3D reconstruction, Bemporad (2009) estimated the expansion factor in three dimensions of an erupting prominence, establishing that it was non-isotropic. In addition, he pointed out that the observed structure was more similar to a 2D ribbon-like feature, without the twisted morphology and untwisting motion usually expected for an erupting flux rope. The measurement of the magnetic field of this evolving structure could confirm and complete these inferred properties.

Gosain et al. (2009) used the STEREO data to follow the evolution of an erupting filament, and establish its high inclination with respect to the solar vertical (about 43°) and its altitude (about 100 Mm) before its ejection. They also observed untwisting motion as the filament rose in the corona.

Among the new techniques developed for stereoscopic purposes we mention the recent effort by Artzner et al. (2010*). Their “difference method” applied to the 304 Å EUVI filtergrams data uses differences between the images taken by the two satellites. In this way the background emission is canceled out, while the filament image is enhanced (Figure 38*).

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Figure 38: The different steps needed to obtain the difference image by Artzner et al. (2010*) and isolate the prominence from the rest of the data. The top panel shows the STEREO-A (right) and STEREO-B (left) images in epipolar orientation after the raw images have been centered and re-scaled. The middle panel shows the STEREO-B image projected onto the STEREO-A view (left) and the STEREO-A view (right). The bottom panel shows the difference of the two images in the middle panel. Image reproduced by permission from Artzner et al. (2010), copyright by Springer.

The 3D reconstruction of prominence eruptions is also useful for understanding the evolving helicity of the prominence environment. The reconstruction performed by Bemporad et al. (2011) allowed the authors to clearly follow the rotational and twisting dynamics of the filament during the eruption, and to compare it with the overlying coronal arcade before and after the eruption. As found by other authors (see the previous paragraph and, e.g., Green et al., 2007), they suggested that the magnetic energy for the eruption was stored in the nonpotentiality of the magnetic environment surrounding the filament and not in the filament itself, and that during the eruption the filament followed the global dynamics of the system. This confirms most of the models in which the energy for eruption builds up throughout the whole filament channel.

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