2 Large-Scale Properties

This section is dedicated to the general properties of prominences and their environment. With this term we indicate the photosphere and the chromosphere under the prominence, and the corona above it. For a complete understanding of this environment, ideally all these layers should be observed at the same time. However, because of the spatial extent and vastly differing temperatures of the regions involved, multi-wavelength and multi-resolution observations are needed. Unfortunately, they are not always available. Hence, the global structure and equilibrium of the prominence environment remain open questions.

2.1 Morphology

The right panel of Figures 2* and 4* show examples of filaments lying in different parts of the Sun. These locations have a common property: a filament is invariably found to overlay the neutral line (PIL) separating concentrations of photospheric magnetic fields having opposite vertical polarities.

At the same location, surrounding the PIL, is the filament channel, which extends up to the corona. The filament channel is thought to be the magnetic structure within which lies the filament plasma. It can be recognized through the properties of the chromospheric fibrils, which are almost aligned with the underlying neutral line as shown in Figure 7*. This does not happen elsewhere outside the channel, where it is possible to find fibrils crossing the PIL.

In most of the filaments observed in Hα and He ii away from the disk center, it is possible to identify a narrow upper dark spine running horizontally along the axis of the structure. Depending upon the direction of the line of sight, it is possible to see barbs, or lateral plasma extensions (see Figures 5* and 6*). The information coming from the high–spatial-resolution images of SDO/AIA and the Doppler motion measurements (EUV and Hα) inside barbs suggest that their magnetic field is inclined with respect to the local vertical to the Sun’s surface. Some barbs seem to reach down into the chromosphere and possibly into the photosphere, as in the case of, e.g., Schmieder et al. (2013) (and probably even deeper into its magnetic field), even though observationally this aspect is still controversial (e.g., Li and Zhang, 2013). For their role in the formation and stability of filaments, barbs and filament ends are important components that need to be characterized in great detail.

Close views of filaments show a fine dark structure made of long thin loops or threads also almost aligned along the prominence spine (the angle is about 20° – 35°), as shown in Figures 4* and 5*. Both chromospheric fibrils and filament threads are considered to trace the magnetic field, so that they suggest a highly-sheared magnetic field (Martin, 1998b*) almost aligned with the neutral line in the photosphere. This is true both for the main body and barbs in filaments. The highly-sheared magnetic configuration appears to have a role in prominence eruptions (see Section 5).

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Figure 5: Hα high resolution image of a filament taken on 6 October 2004 by the Dutch Open Telescope (DOT). The green circles marks the barbs.
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Figure 6: SDO/AIA 171 Å detail of a filament observed on 11 November 2011. Barbs are clearly visible in absorption at the limb and on disk. Courtesy of NASA/SDO/AIA.

Not all filament channels are filled with filament plasma, even though the observation of empty filament channels is quite rare. In certain cases only segments of a filament are seen along the same PIL. This may be the case for a partially-filled filament and/or for different transparency of the medium. To understand where and why filaments form are among the key goals of the study of prominences. What we can tell is that an empty filament channel is part of the filament formation process (see details in Section 4). Time sequences in Hα obtained during the filament formation show that fibril orientations change in time; only when they become almost aligned with the PIL does the filament channel start to be filled up with material (Martin, 1998b*). An example of an apparently-empty filament channel and its aligned fibrils is given in Figure 7*.

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Figure 7: Detail of a complex active region in Hα showing an empty filament channel (between the white arrows) with aligned fibrils along the PIL. Image reproduced by permission from Martin (1998b*), copyright by Springer.

Filaments seen at the limb show different properties. They may look like structures made of fine vertical threads, like that shown in Figure 22*, or of horizontal threads as in Figure 8*. They may also appear to have more complex morphologies, including arcades, like that shown in Figure 9*. Some of these differences may be attributed to different viewing angles at the times the images were recorded. For example, in some cases it is possible that we see just the final end of the structure (the barb), which is compatible with seeing mainly vertical substructures. In contrast an observation of a prominence perpendicular to its spine often reveals almost horizontal substructuring. However, it is not always easy to make this distinction and, if this fine structure traces the fine-scale magnetic morphology, the vertical structuring is not apparently compatible with the horizontal magnetic field deduced from on-disk observations. Possible solutions to this paradox come from measurements of the magnetic field, as discussed in Section 2.4. However, clear differences in the fine-scale morphology exist for which we still do not have a complete explanation.

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Figure 8: High-resolution image on the solar limb obtained with Hinode/SOT Ca ii H-line 3968 Å. Image reproduced by permission from Okamoto et al. (2007*); copyright by AAAS.
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Figure 9: Hα line center image of a prominence at the West limb on 22 November 1995. Credits: Big Bear Solar Observatory/New Jersey Institute of Technology.

If the morphology of a filament marks its magnetic skeleton, interesting information may be deduced during the liftoff of an erupting prominence. In this case, prominences are stretched and isolated from their environment, giving a better view of their structure. Some observations of erupting prominences reveal the presence of a helical structure (see Figures 10Watch/download Movie and 33*), even though it is not clear yet whether this is a consequence of the eruption or whether it was already present during the quiescent period. In fact, there is also evidence for on-disk filaments with a similar structure, which supports the flux rope models. Further details on prominence eruptions are given in Section 5.4.

Figure 10: mov-Movie (12201 KB) One of the first prominence eruptions observed by SDO/AIA. Movie taken from External Link Courtesy of SDO (NASA) and the AIA consortium.

2.2 Plasma parameters

One of the main steps in the comprehension of a prominence is determining its plasma parameters, such as electron density, ionization degree, and temperature. These are obtained through plasma diagnostics of the emitted and/or absorbed radiation, possibly in combination with modeling.

Hydrogen and helium spectral lines and continua emissions are the main tools used for the diagnostics of the cooler prominence core. The outer prominence’s envelope, the PCTR, is hotter so it is instead mainly investigated by the observation of transition-region spectral lines: that is, the emission from heavier elements, but still in low ionization states.

It is important to point out the difficulties in the interpretation of the prominence spectrum, which is produced by both optically-thin and thick plasma. Optically-thick lines such as H Lyα return information on the plasma at different depths along the line of sight as we move from their center toward the wings. Conversely, for optically-thin lines, we deduce information from the integrated emission along the full depth of the emitting plasma (even though the intensity scales as the square of the electron density, for collisionally-excited transitions).

It is extremely important to study both core and PCTR environments, in quiet and dynamic conditions, in order to understand the occurrence of prominences, their stability, and destabilization. However, there is a difficulty when trying to combine the results from these two regions, which are obtained using data with different spatial and temporal resolutions. We recall that the EUV instruments used to observe the PCTR have a spatial resolution about a factor 10 lower than those used to observe the cool core at longer wavelengths. Therefore, the PCTR is observationally much less detailed, and so are the plasma parameters derived in this region.

A recent review (Labrosse et al., 2010b*) describes all the different diagnostic methodologies, together with a set of results, for both prominence layers. Here we present a summary of observational results with some updates.

2.2.1 Prominence core

The prominence core is dense, cool and mostly made of optically-thick plasma. The radiation process is very complex here because the plasma completely or partially absorbs the incoming and auto-produced radiation, and should be diagnosed through solving a non-LTE (non-Local Thermodynamic Equilibrium) radiative transfer problem. Forward methods compare the observed spectrum to a synthetic one obtained from a parameterized prominence model, and solve the associated set of equations for the transfer of radiation through the medium (Labrosse et al., 2010b*). Indirect measurements of plasma parameters can also be obtained using, for example, prominence seismology.

The simplest model used to solve the radiative transfer problem is the 1D slab, whose capabilities in deriving important general properties of prominences have been quite-well analyzed. A clear example of the use of this model is found in Gouttebroze et al. (1993*). These authors analyzed 140 models characterized by different prominence conditions to derive the behavior of the hydrogen emission, and provide its spectral signatures together with the tools for their interpretation.

Starting from the 1D models, 2D slab and multi-thread prominence models have been developed. Unless a multi-thread model is applied, the prominence is considered as a monolithic body, often static, in contrast to what observations show. Besides this approximation, the results revealed the models to be quite good in reproducing the general properties of the structure. Interested readers may refer to Labrosse et al. (2010b*) and references therein for further details on the subject.

The prominence core plasma can also be optically thin to some radiations, such as those produced by less abundant elements. In this case the plasma parameters can be derived without solving the radiative transfer analysis. However, the emitted spectral lines should fall at wavelengths far from the continuum emissions of the most abundant elements H or He (λ > 912 Å), these being the most significant absorbers. Next we review the physical parameters of the prominence core inferred via different techniques.

  • Electron Density. Various methods are used to derive the electron density, for example: broadening of lines due to Stark effect, depolarization by the Hanle effect, or comparing lines with continuum emission. Most of these methods yield a density value between ∼ 109 and 1011 cm–3. From recent observation, Lites et al. (2010*) inferred a few times 1015 cm–3 in an active region filament using a magnetic filling factor deduced from observations in the filament channel and a density deduced in the photosphere. If the hypotheses used in this work are correct, this is one of the highest values reported in the literature. This will have important bearings on the study of the equilibrium conditions of these structures, implying that the stability conditions must be valid over a large range of density values. This should be an element for constraining models.
  • Temperature. The temperature in the prominence core can be derived from H-Ly continuum observations (Parenti et al., 2005*) [the color temperature of which is representative of the electron temperature (Gouttebroze et al., 1993)], or by measuring the kinetic temperature from the width of spectral lines (e.g., Stellmacher et al., 2003*). Figure 11* shows an example of the application of the first method over about 100 Å of a SOHO/SUMER spectrum of a prominence observed on 8 October 1999 (Parenti et al., 2004*). By fitting these data it was possible to derive, under given approximations, an electron temperature of about 8200 K. The values reported in the literature, as derived applying different methods, give a range of values between 7500 and 9000 K.
  • Gas pressure. The pressure in prominences is in the range 0.02 – 1 dyne cm–2. This was already reported in the Hvar reference atmosphere (Engvold et al., 1990b), and has been confirmed by the latest results (Labrosse et al., 2010b*).
  • Mass. Determining the prominence mass is also required to properly model the prominence equilibrium. For this purpose, it is necessary to know the geometrical thickness (and the whole volume) of the structure as well as the ionization degree. There are not many studies aimed at determining the ionization degree. The range of values listed by the Hvar reference atmosphere for hydrogen is still valid (0.2 < N(H+)/N(H0< 0.9, Ruždjak and Tandberg-Hanssen, 1990), while for helium the range is 0.1 < N(He0)/N(H0< 0.2 (Del Zanna et al., 2004, see also Labrosse et al., 2010b* for more details on these results). The typical value for the solar corona is N(He0)/N(H0≈ 0.9 (Feldman, 1992), while the hydrogen is completely ionized above ≈ 3 × 104 K (Mazzotta et al., 1998).

Currently it is acknowledged that the same filament appears wider in the EUV and radio wavelengths than in Hα (e.g., Heinzel et al., 2001, 2003; Schmieder et al., 2003, 2004; Marqué, 2004*; Schwartz et al., 2004, 2006; Vial et al., 2012*). This extra width is called the EUV filament extension. To illustrate, Figure 12* shows the intensity of four different lines (Hα, H Lyα, He ii 304 Å and Fe ix 171 Å) along a cut perpendicular to the filament axis shown in Figure 4* (right panel) and observed on 14 June 2002.

This different appearance of the filament width may be due to the different absorption properties of the H Lyman and Helium continua with respect to the Hα line, to the different amounts of density or filling factor in the two regions, and/or to the lack of high transition region and coronal emission of the prominence. As a consequence, in order to estimate the whole mass of the filament we need to investigate both wavelength regions. In this respect, by comparing the amount of absorption in the EUV lines on top of the H continuum and estimating the Hα absorption, it was established that the optical thickness at the Lyman continuum head (τ912) can be one order of magnitude greater than that of the Hα line center for the same column density. This is also confirmed by non-LTE radiative transfer modeling of filaments by Anzer and Heinzel (2005). In particular, there may be semi-transparent material that is thin to the Hα wavelength and thick to the H Lyman continuum (meaning that the filament is still visible in the continuum emission but not in the Hα images). This is also one reason why the same filament seen in the H Lyman continuum appears bigger than in Hα images. Taking into account this property, the prominence mass established through H and He measurements has been given values in the range of 1014 – 2 × 1015 g (Labrosse et al., 2010b*). Lower values, down to 5 × 1012 g, can be obtained for the neutral hydrogen mass using simple geometrical considerations. A recent multi-wavelength study by Heinzel et al. (2008) established the hydrogen column density of 1 – 5 × 1019 cm–2.

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Figure 11: Radiance of the H i Ly continuum as a function of wavelength in a prominence observed in 1999. The crosses indicate data corrected for temporal variations. The solid curve represents the fit to the corrected data. Image reproduced by permission from Parenti et al. (2005*); copyright by ESO.
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Figure 12: Line intensity profiles along a cut perpendicular to a filament axis observed on 14 June 2002 on the disk and shown in Figure 4* right. Image reproduced by permission from Vial et al. (2012*); copyright by ESO.

2.2.2 Prominence corona transition region

The PCTR emission falls mainly in the UV-EUV and is produced mostly by optically thin plasma. Some of the techniques mentioned before can also be applied in this case. In addition, when multiple spectral lines are available, density, temperature, and Differential Emission Measure (DEM) can be derived.

Differential Emission Measure.
Even if the PCTR accounts for only a small fraction of the mass of the whole cool prominence body, the PCTR demands particular attention because of its interface role. To address this need, there were a few early attempts to build spectral atlases of these structures with the aim of providing as many identified emission lines as possible (e.g., Mariska et al., 1979; Moe et al., 1979; Engvold et al., 1990a). These provide the fingerprints of the structure and the physical parameters obtained through the inversion of this large quantity of lines have low uncertainties. To our knowledge, the most complete spectral atlas available includes more than 400 line profiles in the range 800 – 1250 Å observed with the SOHO/SUMER spectrometer (Parenti et al., 2004, 2005). In addition, the authors built a similar atlas for the quiet Sun observed on the same day, providing additional material for comparison with quiet solar conditions. In terms of prominence plasma parameters, at a given time ideally we would like to determine the plasma distribution in space and temperature. At present, the best that can be done is to try to invert a large set of emission lines to derive the Differential Emission Measure (DEM). This gives the distribution of plasma in temperature, along the line of sight, under the assumption of a constant density. The importance of a large set of lines comes from the need to sample the plasma well enough in temperature to localize the minimum of the DEM, generally found around log T(K) = 5 – 5.3, and the DEM gradient. These quantities, among others, are elements used to test the energy balance in prominences. The DEMs for two different prominences are plotted in the top half of Figure 13*. These are obtained from the SUMER spectral atlas of Parenti and Vial (2007*), and from a similar observing program run on a prominence in June 2004 (Gunár et al., 2011*). The 1999 prominence is estimated to be denser than the 2004 prominence, but what is mostly different is the gradient of the DEM at low transition region temperatures, with the 2004’s DEM being flatter (the coronal DEM is discussed below). It would be interesting to further investigate whether this difference is systematic in prominences or not. At present very few DEMs are available in the literature. When compared to the quiet sun (QS) DEM on the lower panel of Figure 13*, we notice the following differences: lower DEM values at transition region temperatures and a minimum located at lower temperatures. Beside the large amount of lines used to build the curves in Figure 13*, we note that the interval log T(K) = 5 – 5.5 is less well sampled for the QS and prominences alike. This is a general problem with solar observations of the transition region. Different results have been obtained in this context for other prominences (Kucera and Landi, 2006*, 2008*) and the issue is not yet solved. The precision of the DEM distribution in this region could be improved by data in other wavebands obtained with multi-instrument and co-temporal observations. ln addition, we point out that often a given instrument can provide co-temporal data only in a limited range of wavelengths, such as in the case of SUMER. A full spectrum can be covered in up to 1 – 2 hours, depending on the exposure times. This means that the DEMs are built assuming the absence of structural changes in prominences. Therefore, we cannot rule out the possibility that the dynamic nature of these structures has a role in the different results obtained for the DEMs.

It is generally believed that a prominence does not emit at coronal temperatures, and the coronal part of the DEMs of Figure 13* is the result of the foreground and background emission of the off-limb corona. At the same time, the maximum temperature of emission of prominences is still unknown. This is another open issue in prominence physics, whose solution requires a detailed investigation of the emission of the environment around these structures. The difficulties of this task are discussed in Kucera and Landi (2006). Interestingly, the high S/N imagers of SDO/AIA are revealing a faint emission in prominences in the 171 Å channel, for example. The sensitivity of this channel peaks in the Fe ix line, which has its maximum emissivity at 0.8 MK. This would suggest that the PCTR contains more material at hot temperatures than previously thought. Preliminary results on this issue indicate that the prominence emission above 3 × 104 K, which is measured by this channel, is indeed dominated by the Fe ix, confirming that more mass is present at this temperature than expected (Parenti et al., 2012*).

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Figure 13: Differential Emission Measure for two prominences (top) and the QS (bottom). The prominence DEMs were inferred with 20% uncertainty using 1999 (solid line, Parenti and Vial, 2007*) and 2004 (dashed-line, Gunár et al., 2011*) SOHO/SUMER data. The dashed-line on the QS plot is the QS DEM available in the CHIANTI database. Bottom image reproduced by permission from Parenti and Vial (2007*), copyright by ESO.

Electron density and pressure.
A summary of the most-recent electron-density values derived from EUV observations of prominences can be found in Table 4 of Labrosse et al. (2010b*). Most of these results are in the range from ∼ 109 to a few 1010 cm–3 with a few cases over 1011 cm–3. Similar to measurements in the core, in the PCTR we obtain quite a large range of densities and, as expected, they tend to be lower than in the core. This range of values not only includes the variation of density from one prominence to another (although there is no special difference between active and quiescent prominences), but also variations inside the same prominence. This is not surprising considering that these measurements have always been made with limited spatial resolution at or above the known prominence fine scale. As will be discussed in Section 3, in fact, prominences are composed of dynamic fine-scale components, smaller than the resolution of the EUV instruments used to observe the PCTR.

2.3 The coronal cavity

Observations at the limb in the optical, EUV and soft X-ray bands generally reveal variation of the flux with height, which traces a well-defined magnetic-field structure extending above prominences. A darker area, the coronal cavity, extends above and around the prominence up to about 0.6 R ⊙ from the solar surface, and it is overlaid by a faint but visible loop system, the arcade (Figure 14*). At present we recognize that the coronal cavity is how the coronal portion of the filament channel appears at the limb. Over larger scales the arcade is generally part of a streamer structure. A historical review of the existing literature on cavity studies can be found in Gibson et al. (2006, 2010*), while in the following we mention a few more recent results.

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Figure 14: Small cavity (left) and a large cavity (right) imaged in white-light polarization brightness from the MLSO Mk4 coronameter (top), and by the EIT telescope onboard the SOHO satellite, at 304 Å (middle) and 195 Å (bottom). In the 304 Å images, the prominences lying within the cavities are clearly visible. The small cavity is visible at 195 Å while the large cavity does not show up at this wavelength. Image reproduced by permission from Fuller and Gibson (2009*); copyright by AAS.

Morphologically cavities may have different appearances, due both to different properties in different cavities and/or to line-of-sight effects. An example of the different appearance of cavities is illustrated in Figure 14*, which shows the SW quadrant of the Sun where two prominences have been observed at three different wavebands. The prominences are seen in emission at chromospheric temperatures (middle panels) in the EUV. The top panels show the corresponding dark cavity observed in white light (WL). While the two prominences have a similar apparent dimension in the plane of the sky, the two cavities differ in size. The bottom panel shows the same area in EUV at coronal temperatures. While the small cavity is seen in absorption in the left panel, a large cavity is not visible on the right. In addition, their shape often suggests an elliptical cross section (Fuller and Gibson, 2009). More systematic observations of these structures reveal that a coronal cavity’s properties change as a function of the solar cycle, being larger and less dark during the maximum of the cycle.

WL images generally show a decrease of intensity of about 25 to 50% (e.g., Fuller et al., 2008*) in the coronal cavity, with respect to the surrounding arcade and streamer. This has been interpreted as a reduced density and has also been confirmed by radio measurements (e.g., Marqué, 2004). The density variation with height, however, results to be flatter in the cavity than in the arcade.

Contrary to the traditional dark aspect of coronal cavities, the recent results from eclipse data by Habbal et al. (2010*) have documented a coronal cavity that is not darkened in WL images. This is illustrated in Figure 15*, which also shows a bright cavity at coronal temperature suggesting the high temperature of this region. Analyzing cases of different prominence projections with respect to the plane of the sky, these authors suggest that the darkening appears mostly when the prominence axis is along the east-west direction. This may be caused by the longer line-of-sight (LOS) integration path in a reduced-density environment with respect to a prominence (and then a cavity) located along the north-south direction. This aspect has been further investigated by Gibson et al. (2010*) using a forward model of a coronal cavity.

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Figure 15: Top: detail of the 2008 eclipse studied by Habbal et al. (2010*). Second and third rows: close-up views of the prominence regions in Fe x, Fe xi, Fe xiii, and Fe xiv, 7878.6 Å continuum (the limb is oriented horizontally, at the bottom of the image). Bottom row: spectral line intensities normalized to their corresponding maximum values (y-axis) vs. P.A. (x-axis), at 1.05, 1.15, and 1.3R ⊙, with green = Fe x, dashed-green = Fe xi, red = Fe xiii, dashed-red = Fe xiv, and black = 7878.6 Å continuum. The horizontal dashed lines correspond to the radial distances, where the normalized emission-line intensities are plotted. Image reproduced by permission from Habbal et al. (2010*), copyright by AAS.

In terms of thermodynamic parameters, as for prominences, the coronal environment of the cavity is still only partially understood. The faint emission of the cavities makes their study difficult and probably accounts for the variability of the measured parameters. Because of the lack of signal in EUV and X-ray, the data inversion for cavities is mostly limited to WL data, from which we may obtain the electron density. The electron temperature and magnetic structures are generally deduced assuming a hydrostatic model, even though recent efforts have been able to abandon this assumption (Kucera et al., 2012*).

Similar to Habbal et al. (2010*), several recent studies converge in finding the cavity as hot as, or hotter than, the surrounding streamer. The EUV measurements by Vásquez et al. (2009) obtained using tomography and STEREO/EUVI data, support this result together with a broad temperature distribution. Habbal et al. (2010*) assigned a value of at least 2 × 106 K to one of the EUV cavities, bright in the hottest available filter (Fe XIV). Slightly lower values were found by Hudson et al. (1999*) and Hudson and Schwenn (2000) using soft X-ray Yohkoh SXT data in the core of the cavity and by Kucera et al. (2012*) using Hinode/EIS data. In some cases measurements imply a temperature substructure that might be due to the presence of fine structure in the magnetic field that cannot be observed directly (Kucera et al., 2012).

Even though cavities have been observed for decades, only recently has it been possible to observe their fine scale dynamics. For example, Schmit et al. (2009*) investigated prominences using the EUV and infrared spectroscopic information of Hinode/EIS (Fe xii 195 Å) and CoMP (Coronal Multichannel Polarimeter).8 For the first time, they detected coherent Doppler velocity structures (several tens of Mm) of 5 – 10 km s–1 in the cavity, along the prominence axis in both directions and lasting for about 1 h (see Figure 16*), whose origin is still unknown. Wang and Stenborg (2010) have instead detected spinning motions in cavities from observations in the Fe xii 195 Å line with STEREO/EUVI. This transverse motion is about 5 – 10 km s–1 in the plane of the sky and can last for 2 – 3 days. Interestingly it is of similar amplitude as the axial motion measured inside prominences (see Section 3.1.2). These authors noticed that the spinning direction is associated with the asymmetry of the magnetic-polarity concentration at one side of the PIL, and is directed from the strongest concentration to the weakest. They interpret it in terms of siphon flow, which has already been postulated in the presence of asymmetric heating (due in this case to the asymmetric magnetic flux) at loop footpoints. In this picture, the flow highlights the magnetic helical flux rope of the cavity. Indeed, several authors interpret the coronal cavity (and the filament channel) as the location of a twisted magnetic structure (the flux rope) embedding the prominence (e.g., Tandberg-Hanssen, 1995*; Lites, 2005*; Ciaravella et al., 2000; Gibson et al., 2010; Habbal et al., 2010, and references therein).

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Figure 16: (a) Image of a cavity and its prominence observed with Hinode/EIS Fe xii 195 Å. The green contour marks the edge of the prominence as seen from SOHO/EIT 304 Å. The blue and red contours show ± 7 km s–1 LOS velocity. (b) LOS velocity in Fe xii 195 Å. Image reproduced by permission from Schmit et al. (2009); copyright by AAS.

Another question is the origin of hot plasma in the cavity. We need to better characterize it and understand whether it can really be hotter than the surrounding streamer. Some continuous heating mechanism should be at work if the plasma is heated locally. Does it originate at the solar surface or within the cavity?

A few explanations have been proposed. For example, Hudson et al. (1999) suggested the presence of hot filamentary plasma along the magnetic structure, which supports and envelops the cool prominence material. This is consistent with the Parenti et al. (2012) interpretion of AIA/SDO prominence images (see also Section 2.2.2). The origin of this hot plasma was suggested to be chromospheric evaporation due to asymmetric heating localized at the footpoints, as discussed in Antiochos and Klimchuk (1991*). Further explorations of this scenario have reproduced the dynamic filamentary mixture of hot and cool plasmas in prominences and surrounding cavities. For example, Luna et al. (2012b*) recently investigated prominence and cavity formation in a 3D model of a sheared magnetic arcade, heated asymmetrically at the footpoints. Their simulations imply that transition region and coronal temperature plasmas surround every prominence thread, both parallel and transverse to the field.

One of the first prominence investigations with the SDO/AIA imagers (Berger et al., 2011*), in combination with earlier high-resolution studies with Hinode/SOT (Berger et al., 2008*, 2010*), has motivated an alternative concept: a cycle driven by a magneto-thermal convective instability, in which hot plasma escaping from the solar surface passes through prominences, altering their structure and supplying the corona cavity with hot plasma. This promising idea is still under investigation (Berger et al., 2012; Liu et al., 2012; Low et al., 2012b,a), although it is not clear yet how the hot material is produced at the base of the corona. In addition, this model is based on observations of polar-crown prominences, which generally have significant vertical motions and structures (see also Section 3). It is still to be seen whether this instability could work in prominences characterized by different structural properties, and whether this instability could provide enough energy to initiate eruptions.

Both sheared arcade and flux rope models can produce the cavity (for example, see Figure 3*). The presence of a detached flux rope can explain the observed flatter density gradient with height (Fuller et al., 2008), the sharp boundary between cavity and streamer, and the spinning motion, but it is unclear at present whether these results are unique to this model. Fan and Gibson (2006) simulations of the evolution of a quasi-static emergence of a flux rope showed the formation of the cavity, where current sheets form at the interfaces of separatrix layers. These regions can potentially be seen as the location for energy dissipation and plasma heating. This configuration and the others mentioned above are all good candidates that need to be further investigated; in particular, the plasma distribution and evolution in these flux-rope filament-channel models have not been modeled with full thermodynamics, as in the Luna et al. (2012b*) and earlier simulations of sheared-arcade filament channels. Clearly, better knowledge of the magnetic configuration is the key to resolving these issues.

2.4 Magnetic field

It is widely believed that, analogous to other solar coronal environments, the magnetic field in prominences plays a dominant role over the plasma (low β environment). However, a recent investigation using a semi-empirical model of massive, well-developed prominences by Anzer and Heinzel (2007*) indicates that the plasma β locally can take values in the range of 0.2 < β < 1.5, from a magnetic-field-dominated plasma to a matter-dominated plasma. Therefore the discussion of the role of the magnetic field in massive prominences can still be debated.

Information on the magnetic field of prominences and their surroundings is sometimes deduced using the morphology and apparent dynamics of the plasma, including the apparent bulk flow thought to highlight the magnetic field lines of the fine structure. However, the reduced emission of prominences at the limb limits the identification of magnetic fine structures, generally well highlighted in other solar structures (e.g., coronal loops). The magnetic field strength and direction can also be inferred from measurements but with large uncertainties, both because of the faint emission of the prominence plasma and the low magnetic-field amplitude. These aspects make the determination of the global magnetic structure in prominences more difficult.

2.4.1 Properties deduced from morphology

The ubiquitous presence of filaments and prominences on the Sun allows the accumulation of a good sample of data, which helps to identify certain morphological characteristics and to classify objects according to common properties. One of these is the magnetic chirality. Optical and EUV-X ray observations have revealed that the magnetic field of the filament environment (filament, filament channel and coronal arcade) follows a handedness rule called chirality, which is described in detail by Martin (1998b*) and illustrated in Figures 17* and 18*. Opposite chilarity appears to be dominant in each hemisphere: negative in the north hemisphere and positive in the south. However, there is no clear observational evidence yet for a change with the cyclic polarity reversal of the polar field (e.g., Pevtsov et al., 2003), as has been predicted by some models (van Ballegooijen et al., 1998).

Once the magnetic polarity of the filament channel is known from photospheric measurements, and we look at the filament from the positive side, the chirality is dextral if the filament axis is directed rightward, and sinistral if directed in the opposite way. When there is no direct access to the magnetic field, the key aspects to identify the “dextral” or “sinistral” chirality in a filament are the extensions of the barbs from the spine: from left-to-right or right-to-left (Figures 17* and 18*).

This chirality rule is also valid for the chromospheric fibrils inside the filament channel. We have already mentioned that they are almost aligned with the PIL and are rooted in small magnetic concentrations having the same polarity as the underlying photospheric network magnetic field. The loop arcade overlying a filament and its barbs also has a chirality pattern, but it is inverse to the filament: the arcade forms an angle which is generally almost perpendicular to the PIL, but the arcade skew is in the opposite sense to the filament chirality (Figure 17*).

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Figure 17: One-to-one chirality relationships for (1) fibril patterns, (2) filament spines and barbs, and (3) overlying arcades of coronal loops, are shown in each column. The patterns in the left column are dominant in the northern hemisphere and those in the right column are dominant in the southern hemisphere. Image reproduced by permission from Martin (1998a), copyright by ASP.
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Figure 18: High-resolution Hα image of a chromospheric filament observed on October 30, 2002 at 14:46 UT. Upper right and lower left corners: Dextral and sinistral patterns. The small filament in the lower right corner exhibits both sinistral and dextral barbs. USAF, ISOON image courtesy D. Neidig; from work of A.A. Pevtsov; External Link

Further information on the magnetic structure may be inferred from the sign of the magnetic helicity,9 which is a measure of how many twists and turns there are in the magnetic field. By definition the helicity is positive for a right-handed twist, and negative if in the opposite direction. In general, the solar magnetic field has negative helicity in the northern hemisphere and positive helicity in the southern hemisphere. Understanding whether this rule also applies to filaments will shed further light on the link they might have with the surrounding magnetic environment, for example, in terms of their formation. Filaments are also associated with coronal eruptions (see Section 5.4), so they may be used to trace the helicity transport in the corona and in the interplanetary medium. Several theoretical models and simulations suggest that magnetic twist arises from the differential rotation, meridional flow, and diffusion, as well as direct transport during flux emergence (e.g., van Ballegooijen et al., 2000; Martens and Zwaan, 2001; Yeates et al., 2008). However, some of these surface effects imply a dependence of the filament chirality on the PIL orientation, which is not supported by the observations (Lim and Chae, 2009*).

Figure 19* shows an example of how the magnetic helicity in the different temperature regimes can be investigated. In this case the author (Chae, 2000*) analyzed the position of the filament threads observed in Hα in comparison with the 171 Å TRACE images and photospheric magnetograms from NSO/KP. In agreement with previous work, he showed that there were threads with positive helicity in sinistral filament channels, and negative helicity for dextral filament channels. The result supports the idea that filaments have the same helicity as their surroundings (including the nearby active region), implying the same origin for these magnetic features. More recently Lim and Chae (2009) arrived at a similar conclusion by finding a strong correlation between filament chirality (deduced from the barbs’ orientation in Hα images of intermediate filaments) and the active region helicity (deduced from SXT images).

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Figure 19: Observed thread crossings in an inverse S-shaped filament. Note that the filament is dextral. (a) High-resolution Hα data and line-of-sight magnetogram taken at BBSO. The positive flux density levels are represented by white contours, and the negative ones by black contours. (b – d) TRACE images at different times. Image reproduced by permission from Chae (2000); copyright by AAS.

2.4.2 Magnetic field inferred from measurements

Spectro-polarimetry is the main method used to infer magnetic field amplitude and direction. The line emission processes in prominences are dominated by the scattering of photospheric radiation, which produces mainly linear polarization due to the anisotropy of the solar incident radiation (e.g., Landi Degl’Innocenti et al., 1987; Bommier et al., 1994*). This polarization is modified if a magnetic field is present. Indeed, the data reveal the presence of fine structure splitting due to the Zeeman effect, and/or depolarization through the Hanle effect (Leroy et al., 1983*; Kim, 1990).

The first polarimetric measurements were carried out by J. L. Leroy using the Pic du Midi10 facility, starting from the 1970s. For a review on the methodology and these first results see Tandberg-Hanssen (1995*) and Paletou et al. (2001*). In recent years the spectro-polarimetry method has evolved significantly. This is due to the development of more sophisticated analysis techniques combined with the availability of higher quality data, provided, for example, by the optical telescopes THEMIS, the near infrared Vacuum Tower Telescope of the Teide Observatory11 (see Figure 20*) and the Dunn Solar Tower12 (DST, NSO/SP, USA). These telescopes measure the flux of the most used lines for polarimetry: the He i D3 multiplet at 5876 Å and 10830 Å, and/or the H i Hα or Hβ (Paletou et al., 2001). In order to infer the magnetic-field amplitude and direction, an inversion method of line flux needs to be used. Because these lines are often optically thick in prominences (e.g., Stellmacher et al., 2003*; Wiehr and Bianda, 2003; Labrosse and Gouttebroze, 2004; Léger and Paletou, 2009) the radiative transfer effects have to be taken into account in their modeling. The difficulties in the inversion of the 10830 Å triplet are discussed by López Ariste and Casini (2002). In general, we need to remember that the inversion method relies on some assumption about the observed atmosphere; hence, we cannot say that we obtain a direct measurement of the magnetic field.

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Figure 20: Stokes parameters from a multi-line inversion of simultaneous and cospatial spectro-polarimetric observations of He i 10830 (left) and D3 (right) in a quiescent prominence, taken with THEMIS on 29 June 2007. Image reproduced by permission from Casini et al. (2009*); copyright by AAS.

A recent detailed study of the He i 5876 and 10830 Å lines by Casini et al. (2009) has shown that the first multiplet is more suitable for diagnosing stronger fields (∼ 10 G and higher), while the He i 10830 Å is most suitable for weaker fields (0.1 – 10 G), due to the different sensitivity of these lines to the Hanle effect.

Modern instruments used to infer magnetic field values have a typical polarimetry sensitivity of about 103 or better, spatial resolution equal to or greater than half an arcsecond, and exposure times of 1 – 2 mins for prominence observations (or coarser spatial resolution for faster exposures, e.g., Stellmacher et al., 2003*). Unfortunately, these are still low resolutions compared to the fine spatial and dynamical scales of filaments.

Measurements of the direction of the magnetic field perpendicular to the line of sight (that is, in the plane of the PIL or filament axis) deduced from polarimetry have an uncertainty called the 180° ambiguity, because linear polarization is independent of the transverse magnetic field direction. In this case the inversion methods for the Stokes parameters cannot distinguish between the two opposite transverse solutions. Several methods have been developed to resolve this ambiguity, and we refer to Metcalf et al. (2006) and López Ariste et al. (2006*) and references therein for details. However, in filaments and filament channels the uncertainty in the measurements of the weak field, the assumptions about the observed atmosphere, and other factors make this issue harder to resolve than in other regions. Observation of prominence chirality may help solve this problem (Martin et al., 2008).

Previous reviews of magnetic field measurements in prominences are found in Démoulin (1998*), Anzer and Heinzel (2007), Paletou (2008), and Mackay et al. (2010*). The most relevant results about the magnetic field include: the strength of the field from polarimetric measurements is between 8 and 10 G on average for quiescent filaments; magnetic field strengths deduced from filament thread oscillations are within this range, even though some strong assumptions (e.g., electron density) weaken the results (Lin et al., 2009*, filament oscillations will be discussed in Section 3.1.3); most recent polarimetric measurements suggest higher values of a few tens of Gauss in localized patches (e.g., Casini et al., 2003*, 2005), while higher values are associated with active region filaments (up to 600 – 700 G from Kuckein et al., 2009). These differences appear to be associated with the recent inclusion of circular polarization into the inversion method. In addition, the instrument spatial resolution has improved substantially from about 5'' in the 1980s and 1990s to less than 1'' today. The low spatial resolution data provide only a lower limit on the magnetic field strength.

The magnetic field strength can also be deduced from radio emission, by using the relation between temperature brightness, circular polarization, and optical depth (Apushkinskij et al., 1990, 1996). The published results are generally in agreement with those from spectro-polarimetry. One recent measurement at the centimeter-wave radio emission (using the Russian RATAN-600 telescope) from a quiescent prominence gives a strength of 100 – 550 G (Golubchina et al., 2008), a value much higher than those inferred from the polarimetric measurements in these regions. This discrepancy highlights once more that much more effort should be made to derive reliable magnetic-field measurements in the corona.

Leroy et al. (1983*) found a slight increase of the magnetic field strength with height in prominences. This is apparently in contrast with Casini et al. (2003), where the authors give the distribution of magnetic field amplitude in patches. This discrepancy as well could be attributed to the much lower spatial resolution data used by Leroy et al. (1983), but further investigation of this issue is needed.

Polarimetric measurements also provide information on the magnetic field direction. In the main filament body the field is directed almost horizontally and lies in a plane that forms an angle of about 35° with the filament axis (e.g., Trujillo Bueno et al., 2002), confirming results from the morphology of fine fibrils and threads observed in the filament channel and in barbs (e.g., Lin et al., 2008*). However, these results are apparently in contradiction with the observation of the fine structure of some quiescent prominences (as those belonging to the polar crown) where vertical threads dominate. We will come back to discuss this point in Section 3.

At the photospheric level, it is often possible to observe small areas of opposite polarity with respect to the dominant one, on either side of the PIL. These minor polarities (also called spurious or parasite) are candidates for the location where filament barbs are rooted, as discussed in Section 2.4.3 below.

2.4.3 Interpretation of the global magnetic configuration

Joining the measurements and observations of morphology and/or dynamics at photospheric and chromospheric levels, it is possible to infer the global configuration of the magnetic structure embedding the filament at different levels.

To this purpose, the 3D reconstruction of prominence morphology is certainly an important constraining element. When we try to infer the 3D morphology of prominences from their emission we are aware that the angle of view, amount of opacity (prominence mass), and temperature of the structure with respect to the temperature sensitivity of the observing instruments all matter. The observed differences among prominences may be partially due to these elements. For example, does a characteristic length, height, and width of prominences exist? We know that these parameters depend on how magnetically active the environment is where the prominences form. While efforts were made in the past to classify these different parameters, the most recent observations reveal a wide spectrum of dimensions with less definite limits between active, intermediate, and quiescent filaments. This may suggest a continuity in spatial scales of the strength and topological complexity of the magnetic field. Clearly, if the 3D morphology of prominences were observable we could better infer the enveloping magnetic field structure (for example, helical vs. sheared arcade) and the forces governing prominence support and stability. Unfortunately, due to the difficulty of isolating the emission of the lower portion of the structure (or the full structure if it does not extend much in height) from the surrounding bright environment, this reconstruction is, as of today, not easily accessible. In addition, we know that magnetic field measurements, which scan the structure in height, are also difficult. On the contrary, interesting results have been achieved in the case of erupting prominences (where the contrast in brightness with the corona is stronger), as discussed in Section 5.5.

Concerning magnetic field measurements, in addition to the simple configuration where the magnetic field crossing the prominence has the same direction as the underlying photospheric field (as in a simple arcade, called normal polarity), polarimetric measurements have revealed the presence of the inverse polarity of the prominence magnetic field in most cases. The inverse term refers to two directions: the magnetic field perpendicular to the filament axis has a direction opposite to that of the overlying prominence arcade; the field direction along the prominence axis is opposite to what is expected from photospheric magnetic-field extrapolations of the underlying bipolar field (that is, directed from the negative photospheric polarity to the positive one, see Figure 21*).

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Figure 21: Detail of the vector photospheric magnetic field in regions around a filament channel observed on 11 December 2007. Arrows of equal length show the orientation of the horizontal component of the magnetic field vector. The field is directed mainly along the filament channel running from lower left to upper right, but over most of its length the field has inverse configuration: the arrows have a component directed from negative polarity (lower right) toward positive polarity (upper left). This is especially true in the naked bald patch region above and to the right of the center of the image where the filament channel is not flanked by strong plage. A naked bald patch region is characterized by an inverse configuration of the transverse field, while the longitudinal component is almost absent. The gray scale ranges from 0 (white) to 2000 Gauss (black). Low-polarization pixels where no inversion technique has been applied appear as white areas. Image reproduced by permission from Lites (2009), copyright by Springer.

Several 2D and 3D models have been proposed that result in these possibly stable configurations, including the above-mentioned sheared arcade (which produces mixed polarity), the flux rope (with normal or inverse polarity), quadrupolar regions (inverse polarity), etc. We suggest consulting Tandberg-Hanssen (1995*) and Mackay et al. (2010*) for details. Some of these configurations are illustrated in a schematic view in Figure 3*. Statistical observational studies reveal that most of the prominences have inverse polarity (between 75% and 90%, Démoulin, 1998*), while normal polarity prominences have a stronger field and are located closer to active regions (e.g., Bommier et al., 1994).

We point out that the identification of inverse or normal polarity in quiescent filaments is very difficult because of the weak magnetic field; for correct measurement, high spatial resolution and instrument sensitivities are required (Lites, 2005), as well as systematic measurements at different heights, particularly above the photosphere.

Affirming the presence of a flux rope in the data is also quite difficult; only in the last decade have a few cases been found, but only near active regions. Erupting filaments (see Section 5.4) often appear to be twisted, as expected if their magnetic envelope were a flux rope (see Schmieder et al., 1985; Vršnak et al., 1988; Wang et al., 1996, for a few early examples and Figures 31* and 10Watch/download Movie). However, this does not imply that such a configuration was already present before eruption, as further discussed in Section 4. In conclusion, it is possible that the results on the distribution of inverse and normal polarities just mentioned might be affected by uncertainties in inferring the magnetic field from observations.

The question arises on how the prominence plasma is supported inside these different magnetic configurations. Most stationary models require the presence of magnetic dips (Figure 3*), which create a vertical Lorentz force that can balance the plasma gravitational force. Dips can be formed by the weight of the filament plasma through deformation of the potential magnetic field, or can be formed independently, before the channel is filled by prominence plasma. However, the latter possibility is better supported by the observations: the time observed to fill a filament channel (days) is generally much longer than the time estimated to deform the field lines under plasma weight (hours, Démoulin, 1998). Observationally a dip is consistent with a slight increase in the magnetic field amplitude with height. Their presence is also suggested by observations as highlighted by dynamical studies at small scales, as discussed in Section 3.1.2. For example, their presence can explain the observed periodic longitudinal oscillations of prominence threads, as shown by Luna et al. (2012a*) and Luna and Karpen (2012*).

Models built to explain signatures of the presence of dips have shown that dips can exist, for example, in sheared arcades, in a flux rope, or in a quadrupolar configuration. The question is, can we find observational signatures to distinguish among these different models? High resolution studies of plasma flows (which could resolve the motion of the plasma filling the dips) could allow to distinguish between flux rope and sheared arcade magnetic configurations by detecting different observational signatures: helical or almost horizontal flows. From the model side, they can indeed reproduce several of the observed properties of dips, but they are not completely mutually exclusive (e.g., Karpen et al., 2003). Indeed, observations suggest that at the present reachable spatial resolution, filaments with different magnetic configurations exist.

As mentioned, additional information on the magnetic field can be obtained from the investigations of barbs and their endpoints. For example, signatures of dips have also been found in barbs, and their endpoints were found to be located in small bipoles of parasitic photospheric polarities (e.g., Zong et al., 2003; López Ariste et al., 2006; Martin, 1998b). These again are signatures of inverse polarity of the magnetic topology (see also Xu et al., 2012). Lin et al. (2005*) inferred the photosphere as the anchorage location of prominences. They found barbs to root in faint magnetic-field regions located in the intergranular part of the photosphere, even if they do not give any indication about the polarity of those regions.

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