1 Introduction

Solar prominences are intriguing structures of the solar atmosphere. As shown in Figures 1* and 4*, they appear as bright arcade-like structures (or long and thin clouds) on or above the solar limb when observed at chromospheric temperatures (3 4 7 × 10 ≤ T ≤ 2 × 10 K). They are not seen, or are seen in absorption, at coronal temperatures (Figure 1*). These observations suggest that these features are made of dense cool chromospheric material immersed into the 1 MK corona. The interface between the two environments is called the Prominence-Corona-Transition Region (PCTR). As a consequence of the low temperature, the prominence core is made of partially-ionized plasma and, because of its high density, it is optically thick to certain wavelengths (e.g., most of the hydrogen and helium resonance lines and continua). For this reason, when prominences are observed on the disk, for example in Hα 6562.8 Å or in He ii 304 Å (which are among the brightest lines produced by the chromosphere) they appear darker than the surrounding quiet Sun (see Figures 1* and 2*, these observations are discussed in Section 1.3). In this case they were historically called filaments. For simplicity, in the present work we will use both terms interchangeably to indicate these features.

Prominence properties may vary over wide ranges. Based on these differences, various classifications of prominences have been made based on their activity, morphology, location and on the properties of the magnetic field in the photosphere below them. A review of their classification can be found in Tandberg-Hanssen (1995*). Updates, when necessary, will be given in the present text.

Prominence properties depend upon the environment where they form, and in particular upon the magnetic field below them: filaments are always found above the neutral lines separating opposite polarities of the photospheric magnetic field (Polarity Inversion Line, PIL). The filament width, length and shape follow and adapt to the extension of the neutral line (even though it does not always cover its full length). The extension of the PIL depends upon the strength and distribution of the local magnetic field, so that filaments and the local magnetic field are closely interrelated. Since PILs can be found almost everywhere on the Sun, so can prominences, particularly during the maximum of the solar cycle. For all these reasons their appearance may vary a lot as shown, for example, in Figure 2*.

However, filaments are primarily found all around the border of polar coronal holes (this location is called the polar crown), between active regions or surrounding them (intermediate filaments) and inside active regions (active filaments). In the first and second cases the filaments are generally longer (depending also on the extension of the nearby active region), wider, extend more in height, and have longer lifetimes (particularly the first class) than in the last case. This is probably due to the more stable nature of the environment where they live. These are called quiescent filaments. The longest filaments can cover almost the solar diameter (transequatorial filament), as shown in Figures 1* and 2* marked by the yellow arrows. This suggests a connection between filaments and the large scale magnetic structure of the Sun.

The regular observations of the full disk made historically on the ground by Hα telescopes and later on from space in the He ii 304 Å channel (e.g., by SOHO/EIT, STEREO/EUVI or SDO/AIA) are sources of data for statistical studies aimed at determining general prominence and filament properties. In order to help the statistical analysis, several automated detection procedures have been developed or are under development (e.g., Bernasconi et al., 2005*; Romeuf et al., 2007; Aboudarham et al., 2008; Wang et al., 2010*; Labrosse et al., 2010a; Joshi et al., 2010; Buchlin et al., 2010). For example, Bernasconi et al. (2005) analyzed almost five years of BBSO Hα filtergrams (19 211 filaments between 2000 and 2005), which revealed that the average filament length varies from about 3 × 104 to about 1.1 × 105 km. Active filaments are the shortest and last from minutes to hours.

The recent EUV prominence statistical study by Wang et al. (2010) confirms these values. These results are obtained from the analysis of 9477 limb prominences in the period 2007 – 2009 (minimum of solar activity during the end of cycle 23 and the beginning of cycle 24) observed with the STEREO/EUVIs 304 Å imager. Their results also show that the long prominences were located mostly between 30° and 60°; 82% of them had a height of about 2.6 × 104 km above the solar surface. However, quiescent prominences can reach much higher altitudes. Their brightness is almost constant, or slightly diminishing with height.

No detailed statistical information is found in the literature on the filament width; however the reported values vary between 103 and 104 km (e.g., Malherbe et al., 1983; Pouget et al., 2006; Labrosse et al., 2010b*).

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Figure 1: A prominence at the limb and its filament part on the disk as seen by SDO/AIA. The off-limb prominence appears bright in the 304 Å passband (left panel, chromospheric emission) while the on-disk part, the filament, appears darker than the surrounding quiet Sun suggesting absorption effects. The filament is still dark in the 171 Å band on the disk image (right panel, coronal emission) and off-limb it is almost invisible, suggesting the lack of emission at these temperatures. Courtesy of NASA/SDO and the AIA science team.
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Figure 2: Left: image of the Sun as seen in the He ii 304 Å channel on SOHO/EIT, on 23 February 2004, revealing three dark filaments (arrows). One of them extends over almost a solar diameter. Right: image of the Sun in Hα on the same day revealing different details of these three structures. Credits: ESA/NASA/EIT; Catania Astrophysical Observatory.

1.1 Motivations for prominence study

Decades of study have shown that it is very difficult to characterize filament properties. Filaments show differences in morphology, lifetime, position on the solar disk, complexity of their magnetic field environments, etc. They are not uniform in shape and show a fine, dynamic structure at the limit of the instrumental resolution. This high variability makes their classification difficult and also results in a wide range of physical conditions deduced from observations that poorly constrain the models of prominence formation and disappearance (Vial, 1998). At the same time, understanding the origin of such variety and attaining better knowledge of these structures and their environment during the different phases of their life can provide valuable information on the physics of the solar atmosphere. In fact, prominences are commonly found in the solar atmosphere, which indicates that it is easy to find favorable conditions for their formation and stability. This tells us that prominences are manifestations of a common physical process found in the solar atmosphere. An overview of the large scale properties of prominences, including the morphological aspects, will be given in Section 2.

The following still-unsettled issues motivate prominence studies:

  • Support and stability. In the quiescent state, prominences are interesting for their puzzling equilibrium condition that allows their mass to be supported in the tenuous corona. Clearly the magnetic field plays a major role, but still we do not have sufficient observational information to identify the different mechanisms and/or magnetic configurations responsible for their stability.

    Figure 3* shows the magnetic topology of the first static prominence model proposed by Kippenhahn and Schlüter (1957). The prominence is represented by a thin sheet of current, perpendicular to the magnetic field. In this configuration the gravity force of the prominence is balanced by the Lorentz force. Since then prominence models have evolved and, as reviewed by Mackay et al. (2010*), they may be divided into two main groups: those in which the prominence mass has a major role in the stability of the magnetic configuration, and those in which its role is not important. Figure 3* also shows the two main magnetic configurations thought to host filament plasma: the sheared arcade (middle panel) connecting opposite magnetic polarities at the side of the neutral line (thick-light lines) overlain by an unsheared arcade (solid-dark lines), and the flux rope (right panel) where the magnetic field has a helical configuration. In all cases the stationary prominence material is found in or nearby the dipped regions (concave-up fields, as in Figure 3* ). Section 2.4 reviews the information on the magnetic field in prominences and their surroundings that led to the development of these models.

  • Mass motions. Mass flows inside a prominence can also have a role in the equilibrium, mass support, and mass refurnishing. As we will see in Section 3, observations reveal a variety of mass motions inside prominences, even though the spatial resolution of instruments may limit their diagnostic capability. However, only a few prominence models include such observed dynamics.
  • Radiative losses. Partially-ionized prominence plasma is an interesting laboratory for testing our knowledge of the radiation-transfer mechanisms in an optically-thick medium. In addition, non-local thermodynamic conditions (NLTE) generally exist in these plasmas. Even if the prominence density is quite high (and so the mean free path of electrons is small), collisions are unable to compete with radiation in populating the energy levels of the atoms, so that the Local Thermodynamic Equilibrium Condition (LTEC) does not hold. The incoming radiation due to the environmental solar emission is, in fact, another element affecting the prominence physical conditions. Proper treatment of radiative transfer in prominences is also important for quantifying the amount of radiative losses: this mechanism acts as a cooling mechanism in the energy equilibrium equation. In quiescent conditions, the prominence radiation is steady, requiring a source of still unknown heating to maintain energy balance in the structure.

    The plasma properties derived from observational measurements are given in Section 2.2. A distinction between the properties of the cool prominence core (Section 2.2.1) and its interface with the corona (Section 2.2.2) is also discussed.

  • Magnetic field. The lack of extensive magnetic-field measurements in prominences limits our knowledge of the physics of the coronal magnetic field and its interaction with the plasma. We often assume that the observed prominence–plasma-emission morphology traces the magnetic field lines. One of the paradoxes in prominence studies is that prominences at the limb can show a vertical fine bright structure, while disk observations and the few magnetic field measurements suggest that the field is almost horizontal. Several ideas have been proposed as solutions, and we review them in Section 3. Solving this conundrum will help to understand the coronal magnetic environment and to identify the dominant physical process in the solar atmosphere.
  • Formation and disappearance. Prominences as a whole can be very stable for a few months, or can be part of large-scale dynamic and energetic events such as flares and coronal mass ejections (see for example van Driel-Gesztelyi and Culhane, 2009*, for a review). These enormous eruptions (about 3 × 1015 g on average, Hudson et al., 2006) perturb the interplanetary medium, and their effects can be seen on Earth. For example, they are the origin of geomagnetic storms, which can affect everyday life through electric blackouts. As modern life is very technology dependent, such solar activity is a concern. It is well known that the technology of satellites for telecommunication and human space activity are very sensitive to solar eruptions. For all these reasons the past decade has seen the development of Space Weather as a new branch of science aimed at forecasting solar activity and its consequences on Earth. For accurate predictions it is essential to ascertain whether some phenomena are systematic precursors of solar eruptions, for example, instabilities in the filament-channel magnetic field. The instability conditions and the following disappearance of prominences are reviewed in Section 5. For further discussions on coronal eruptions and space weather see the reviews by Schwenn (2006*), Benz (2008*), and Chen (2011*) in this journal, and Kunow et al. (2006).
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Figure 3: Left: the first proposed magnetic topology for prominence from Kippenhahn–Schlüter (from Gilbert et al., 2001). Middle: sheared arcade configuration (from DeVore and Antiochos, 2000*). Right: flux rope configuration (from Amari et al., 2003b). Images reproduced by permission, copyright by AAS.

1.2 Historical background

Being visible during eclipses, prominences have been known for a long time. However, only during the eclipse of 1860, when photography was introduced for the first time, were prominences finally recognized as solar features and not an effect of the Earth’s atmosphere (Foukal, 2004). In the following years spectroscopy started to develop; during the eclipse of 1868 the yellow D3 line at 5877 Å was observed for the first time in prominences, later identified as coming from solar He emission (Janssen, 1869; Secchi, 1870). Further advances in prominence knowledge and spectroscopic methods came with the discovery that prominence emission could be observed outside the limb even during daylight (Lockyer, 1868) with coronagraphs. In the 1890s the observations of filaments on the disk with spectroheliographs began, and the first systematic photometric measurements in prominences were made by Schwarzschild in 1906. Lyot developed the coronagraph in 1936, ushering in a new era in which systematic prominence limb observations can be made outside of eclipses (Lyot, 1939). The first studies relating prominences to the solar cycle were published by Bocchino (1933) and Barocas (1939). Further information on historical prominence observations can be found in Tandberg-Hanssen (1974, 1998) and references therein.

Daily images of the Sun from the ground and from space are available today.1 They allow us to fully record solar activity in all structures, including prominences. These observations are presented in the next Section 1.3.

1.3 Prominence observations today

Similar to the rest of the Sun, prominence material is made almost completely of hydrogen and helium. Because of their low temperature, prominences abound in neutral or low ionization charge states. Observations of their emission/absorption aim at inferring their cooler core properties. For these reasons, prominences are best observed in the intense hydrogen and helium Lyman and Balmer lines series. Most often, observations are taken in the strong red visible Balmer Hα line, be it from the ground2 or from space (Hinode/SOT, Tsuneta et al., 2008). This long wavelength allows spatial resolutions of fractions of an arcsecond to be achieved.

When observing on the disk, the Hα line is optically thick most of the time. For this reason, by scanning the line at different wavelengths it is possible to image different plasma layers: the filament is seen at line center, while at about 0.5 Å away from this position we see the chromospheric fibrils below it (see Section 2.1 for details on these structures).

Being dynamic objects, prominences require short exposure times. These could be achieved using bright lines such as the UV H Lyα (1215.67 Å) and Balmer Hα. However, in the UV passband the great difference between the intensity of the H Lyα and the other lines produces a technical challenge when trying to observe all of them with the same instrument. For this reason often the easiest choice is made by sacrificing one or the other wavelength window, unfortunately introducing a restriction in the scientific return from the instrument. For example, this is the case for the spectrometer SOHO/SUMER, which uses an attenuator to partially obscure this line, but accesses the other lines of the H Ly series (at about 1'' spatial resolution). This is a limitation for the diagnostics of prominences. Improvements were obtained later in the SOHO mission with increased interest in this line. New solutions for the operation modes were found, allowing access to the full line profile (e.g., Wilhelm et al., 1995, 2004, 2007; Ebadi et al., 2009; Curdt and Tian, 2010).

In conjunction with its high radiance, the Lyα line is optically thick most of the time, which, similar to Hα, makes its interpretation very difficult. The complexity of this emission has been recently illustrated by Vourlidas et al. (2010*) using data from the VAULT rocketborne telescope.3 VAULT (Korendyke et al., 2001) imaged the chromosphere and the bottom of the transition region in Lyα at the best spatial resolution to date (about 0.15''). An example of a prominence in emission at the limb and a dark filament on the disk are given in Figure 4*, where the very fine scale structure of these objects can be observed. The VAULT data study by Vourlidas et al. (2010*) revealed that the Lyα flux in filaments is highly anisotropic and that it may be optically thin in certain parts. Unfortunately, the line profile was not accessible.

These difficulties explain why the Lyα prominence observations have not been very popular thus far, and only a few instruments have been dedicated to it (for example, see the TRACE4 mission Handy et al., 1999).

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Figure 4: VAULT Lyα images of a prominence at the limb (left) and a filament on the disk (right) observed on June 14, 2002 (see Vourlidas et al., 2010).

The near infrared absorption triplet at 10830 Å and the D3 multiplet 5876 Å of He i are also often used for prominence observations and magnetic field measurements (see Section 2.4). The prominence-corona-transition-region (PCTR) is easily accessible through space observations of the EUV He ii 304 Å resonance line. As mentioned already, several missions provide images of the full Sun at this wavelength (for example, SOHO/EIT5, STEREO/EUVI6 and SDO/AIA7). Spectroscopic measurements of several of the most intense EUV He i and ii lines are made by the SOHO/CDS and Hinode/EIS spectrometers (Harrison et al., 1995; Culhane et al., 2007). Alternatively, the Ca ii H-line 3968.5 Å is another chromospheric line used to observe prominences from space with ground based instruments.

Because of the strong absorption of the H and He continua, filaments on disk appear as dark features when observed in the EUV-FUV below 912 Å, that is in the H Lyman continuum head (504 Å for He i and 228 Å for He ii). This property is used to investigate the opacity of the medium and derive information on the prominence mass and degree of ionization. Furthermore, the dark aspect of filaments at high temperature emission observed in this waveband may come from what is called volume blocking, which means the lack of emission at such high temperatures (Heinzel and Anzer, 2001*; Schwartz et al., 2006*, and reference therein). At wavelengths longer than 912 Å filaments are often not visible on the Sun at chromospheric temperatures, meaning that their plasma is transparent to such wavelengths and that their emission is probably too faint to be measured against the disk emission. Instead, at the limb, at chromospheric temperatures the optically thin emission of a prominence has been estimated to be about one third that of the solar disk. Emissions from different layers of prominences, together with the thermal structuring of the PCTR, are useful for evaluating the radiation losses and energy balance. Despite being addressed by several studies (see for example Del Zanna et al., 2004*; Parenti and Vial, 2007*, and reference therein), these prominence aspects are still poorly understood.

Prominences can also be observed at longer wavelengths. In particular, they are bright at the limb against the ambient corona when observed in the microwave range. On the contrary, they appear in absorption, similar to Hα images, when observed on the disk. Due to the density dependence of the radio wave propagation and the optical thickness of the medium, only the shorter radio waves (λ ≤ 1 cm) can emerge from the deepest (and densest) prominence layers. Longer wavelengths can be used to sample the PCTR (Chiuderi Drago, 1990). It should also be taken into account that, due to the non-uniform geometry of prominences and the density-dependent radio-wave propagation, radio observations at the limb provide information on a different layer of the prominence than sampled in disk observations.

To conclude, we mention the newly launched Interface Region Imaging Spectrograph (IRIS) NASA mission. Its high resolution instruments (about 0.4'') can provide prominence data, such as the Mg ii k&h lines profile. These are among the brightest lines of the solar chromosphere. To date the literature does not yet include results on prominence studies with IRIS, but throughout this review the diagnostics capability of this instrument will be discussed.

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