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Solar prominence
Solar prominence
Solar prominence
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Solar prominence seen in true color during totality of a solar eclipse.

In solar physics, a prominence, sometimes referred to as a filament,[a] is a large plasma and magnetic field structure extending outward from the Sun's surface, often in a loop shape. Prominences are anchored to the Sun's surface in the much brighter photosphere, and extend outwards into the solar corona. While the corona consists of extremely hot plasma, prominences contain much cooler plasma, similar in composition to that of the chromosphere. Like the corona, solar prominences are only visible to the naked eye during a total solar eclipse.

Prominences form over timescales of about a day and may persist in the corona for several weeks or months, looping hundreds of thousands of kilometers into space. Some prominences may give rise to coronal mass ejections. Exact mechanism of prominence generation is an ongoing target of scientific research.

A typical prominence extends over many thousands of kilometers; the largest on record was estimated at over 800,000 km (500,000 mi) long,[2] roughly of solar radius.

History

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The first detailed description of a solar prominence was in 14th-century Laurentian Codex, describing the solar eclipse of May 1, 1185. They were described as "flame-like tongues of live embers".[3][4][5]

Prominences were first photographed during the solar eclipse of July 18, 1860, by Angelo Secchi. From these photographs, altitude, emissivity, and many other important parameters were able to be derived for the first time.[6]

During the solar eclipse of August 18, 1868, spectroscopes were for the first time able to detect the presence of emission lines from prominences. The detection of a hydrogen line confirmed that prominences were gaseous in nature. Pierre Janssen was also able to detect an emission line corresponding to an at the time unknown element now known as helium. The following day, Janssen confirmed his measurements by recording the emission lines from the now unobstructed Sun, a task that had never been done before. Using his new techniques, astronomers were able to study prominences daily.[7]

Classification

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Grayscale H-alpha image of the Solar disk showing quiescent filaments (QF), intermediate filaments (IF), and active region filaments (ARF).

Historically, any feature that was visible extending above the surface of the sun, including solar spicules, coronal loops, and some coronal mass ejections, was considered a solar prominence. Today, due to a better understanding of the diversity of these phenomena, most of these are classified separately, and the word prominence is primarily used to refer to larger and cooler features.[8]

There are a number of different prominence classification schemes in use today. One of the most widely used and basic schemes classifies prominences based on the magnetic environment in which they had formed. There are three classes:

  • Active region prominences, or active region filaments, form within the relatively strong magnetic fields at the centers of active regions. Active region prominences have lifetimes ranging from hours to days and erupt more often than prominences belonging to the other classes.[9] As a consequence of being located within active regions, active region prominences are usually found in low heliographic latitudes.[10][11]
  • Intermediate prominences, or intermediate filaments, form between a weak unipolar plage regions and active regions.
  • Quiescent prominences, or quiescent filaments, form in the weak background magnetic field far from any active regions.[12] Unlike active region prominences, quiescent prominences are relatively stable and can have lifetimes ranging from weeks to months, hence the name quiescent.[9] Quiescent prominences are typically located at high latitudes around what is referred to as the polar crown.[10][11] Additionally, quiescent prominences generally reach much greater heights in the corona than active region prominences.

Active region and quiescent prominences can also be differentiated by their emitted spectra. The spectra of active region prominences is identical to that of the upper chromosphere having strong He II lines but very weak ionized metal lines. On the other hand, the spectra of quiescent prominences is identical to the spectra measured at 1,500 km (930 mi) in the chromosphere with strong H, He I, and ionized metal lines, but weak He II lines.[13]

Morphology

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Filament channels

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Prominences form in magnetic structures known as filament channels where they are thermally shielded from the surrounding corona and supported against gravity. These channels are found in the chromosphere and lower corona above divisions between regions of opposite photospheric magnetic polarity known as polarity inversion lines (PIL).[b] The presence of a filament channel is a necessary condition for the formation of a prominence, but a filament channel can exist without containing a prominence. Multiple prominences may form and erupt from within one filament channel over the channel's lifetime. The magnetic field making up the filament channel is predominantly horizontal, pointing in the same direction on both sides of the PIL (see § Chirality).[14][15][16]

Prominence material does not occupy the entire width of the filament channel; a tunnel-like region less dense than the corona, known as a coronal cavity, occupies the volume between the prominence and the overlying magnetic arcade.[7]

Spines and barbs

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Typical prominences have a narrow structure oriented along the filament channel known as a spine. The spine defines the upper main body of a prominence and is generally in the form of a vertical sheet that diverges towards the photosphere at both ends. Many prominences also have smaller structures referred to as barbs that similarly diverge from the spine towards the chromosphere and photosphere. Spines and barbs are both composed of thin threads that trace the magnetic field similar to chromospheric fibrils.[15]

The cool prominence material that makes up spines and barbs—the prominence core—is surrounded by a prominence-corona transition region (PCTR) where there is a steep temperature gradient. The PCTR is responsible for most of the optical emission of prominences.[7]

H-alpha image of an active region filament showing a spine, two barbs, and chromospheric fibrils oriented parallel to the PIL[14]

Overlying structures

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Main article: Helmet streamer

Above filament channels lie overarching magnetic arcades which can extend from 50,000 to 70,000 km (31,000 to 43,000 mi) into the corona. Above these arcades, the closed coronal magnetic field may extend radially outward, forming what is known as a helmet streamer.[17] These streamers may reach a solar radius or more above the Sun's surface.[7]

Chirality

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Filament channels and their prominence, if present, exhibit chirality. When observed from the side of the filament channel with positive magnetic polarity, the channel is said to be dextral if the horizontal magnetic field is oriented rightward and sinistral if it is oriented leftward. Dextral channels have been found more frequently in the Sun's northern hemisphere and sinistral channels more frequently in the southern hemisphere.

The horizontally oriented magnetic field causes chromospheric fibrils along the filament channel to lie nearly parallel to the PIL and anti-parallel to one another on opposite sides of the PIL. The directions that these fibrils are oriented depend on the chirality of the channel. On the side of the PIL with positive magnetic polarity, dextral channels have fibrils which stream to the right and barbs which bear to the right, whereas sinistral channels have fibrils which stream to the left and barbs which bear to the left. Additionally, the overlying magnetic arcades of dextral channels are left-skewed, and those of sinistral channels are right-skewed.[7]

Formation

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The exact mechanism which leads to the formation of solar prominences is not currently known. Models must be able to explain the formation of the filament channel and its hemisphere-dependent chirality, as well as the origin of the dense plasma that makes up the prominence core.[7]

Eruption

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See also: Coronal mass ejection
A solar prominence erupting. False color ultraviolet image.

Some prominences are ejected from the Sun in what is known as a prominence eruption. These eruptions can have speeds ranging from 600 km/s to more than 1000 km/s.[1] At least 70% of prominence eruptions are associated with an ejection of coronal material into the solar wind known as a coronal mass ejection.[18]

See also

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Explanatory notes

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References

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Further reading

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Revisions and contributorsEdit on WikipediaRead on Wikipedia

Solar prominence

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A solar prominence is a large structure of relatively cool and dense plasma suspended in the Sun's hot corona, typically appearing as bright, loop-shaped or cloud-like features extending from the solar surface when observed against the dark background of space, and as dark, elongated filaments when projected against the brighter solar disk. These structures are composed primarily of partially ionized hydrogen and helium, with temperatures ranging from about 7,000 K to 20,000 K—roughly 100 times cooler than the surrounding corona—and electron densities of 10^{10} to 10^{11} cm^{-3}, making them optically thick in chromospheric spectral lines such as Hα at 6562.8 Å. Anchored to the photosphere along polarity inversion lines in the magnetic field, prominences can span tens to hundreds of thousands of kilometers in length and height, often forming arched loops guided by tangled magnetic field lines that provide the necessary support against gravity. Solar prominences form through complex interactions of the Sun's , typically emerging over about one day as cooler plasma condenses in ropes or along neutral lines, and they can persist for days to several months depending on stability. They are classified into main types: quiescent prominences, which are long-lived and often located at the borders of polar or between active regions; active prominences, shorter-lived and associated with groups in active regions; and occasionally transequatorial prominences, which span across the solar equator. Observations reveal internal dynamics, including flows, oscillations, and threads of cooler material, best captured by space-based telescopes like NASA's and ESA's in wavelengths. Prominences play a crucial role in solar activity, as their eruptions—triggered by magnetic instabilities—can lead to coronal mass ejections that propagate through the , potentially disrupting Earth's , satellites, and power grids via geomagnetic storms. Their visibility during total solar eclipses or with coronagraphs has historically aided in mapping the corona, while modern multi-wavelength imaging continues to refine models of their magnetic support and evolution, underscoring their importance in understanding .

Fundamentals

Definition and Characteristics

A solar prominence is a dense, relatively cool structure of plasma suspended within the hot solar corona, often manifesting as bright, loop-like or arc-shaped features when observed at the Sun's limb in (H-alpha) emission, while the same structures appear as dark, elongated filaments when viewed against the bright solar disk due to absorption of background radiation. These structures are primarily composed of partially ionized and plasma, along with trace heavier elements, and are threaded and supported by strong that isolate them from the surrounding environment. Key physical characteristics of solar prominences include core temperatures ranging from approximately 7,500 to 9,000 K, significantly cooler than the million-kelvin corona, with electron densities in the core typically between 10^9 and 10^11 particles per cubic centimeter—about 10 to 100 times denser than the ambient coronal plasma. Their sizes vary, but average lengths span 30,000 to 110,000 km, widths are on the order of 1,000 to 10,000 km, and heights can reach up to 26,000 km, though quiescent prominences may extend higher. Lifetimes differ by type, with quiescent prominences persisting for days to months, while more active ones may last only hours before evolving or erupting. In contrast to the surrounding corona, which exhibits temperatures exceeding 1 million and densities around 10^8 particles per cubic centimeter with full , prominences maintain their lower temperature and higher through magnetic confinement, creating a thin prominence-corona transition region where properties gradually change. This isolation prevents rapid heating by , allowing the prominence plasma to remain in a cooler, denser state despite being embedded in the tenuous, hot coronal medium.

Observation Methods

Ground-based observations of solar prominences primarily rely on spectroheliographs and narrowband H-alpha filters to capture their structure. Spectroheliographs, developed in the late , scan the solar disk in the H-alpha line to produce monochromatic images where prominences appear as bright features against the dark limb or as dark filaments when viewed on the disk. Modern ground-based telescopes, such as those at the Big Bear Solar Observatory, use H-alpha filters to achieve resolutions down to 0.3 arcseconds, resolving fine threads and detecting plasma flows within prominences. These observations are limited by Earth's atmosphere, which introduces seeing distortions, but in facilities like the Dutch Open Telescope enhance image quality for detailed morphology studies. Space-based observatories provide uninterrupted, high-cadence views free from atmospheric interference, enabling multi-wavelength and of prominences. The (SOHO) employs the Large Angle and Spectrometric Coronagraph (LASCO) for white-light observations of prominence eruptions in the corona and the Solar Ultraviolet Measurements of Emitted Radiation (SUMER) instrument for ultraviolet (UV) , capturing line profiles at 1 arcsecond resolution to diagnose plasma properties. The (SDO) Atmospheric Assembly (AIA) images prominences in (EUV) bands such as 304 Å and 171 Å at 12-second cadences, revealing multi-temperature structures and faint coronal emissions associated with prominences. The Hinode X-Ray Telescope (XRT) observes prominences in soft X-rays, highlighting hot coronal cavities and threads invisible in cooler wavelengths. The Interface Region Spectrograph (IRIS) provides high-resolution (0.4 arcsecond) UV in lines like Mg II k&h, enabling plasma diagnostics in the prominence-chromosphere interface. More recent missions, such as the ESA/ launched in 2020, contribute advanced observations through its Imager (EUI) for high-resolution EUV of prominence structures and dynamics, and the of the Coronal Environment () for EUV , capturing detailed plasma properties and eruptions as of 2025. Key spectral lines are essential for diagnosing prominence plasma conditions, with observations targeting optically thick emissions from cool, dense regions. The H-alpha line at 656.3 nm is widely used for imaging the prominence core, as it originates from hydrogen atoms at temperatures around 10,000 K and highlights absorption against the disk or emission at the limb. The He I D3 line at 587.6 nm probes cooler prominence regions and is sensitive to magnetic fields and transition region heating. In the EUV, the line at 121.6 nm, observed by instruments like , reveals optically thick core structures and links to the prominence-corona transition region at resolutions up to 0.15 arcseconds. Doppler imaging techniques utilize line-of-sight shifts in lines to map plasma motions within prominences. By analyzing Doppler shifts in H-alpha, observers detect counter-streaming flows with velocities of 2–30 km/s, revealing dynamic threads and barbs. Three-dimensional reconstruction methods, such as from multi-viewpoint data (e.g., from or IRIS), infer prominence geometry and helicity by combining EUV spectra and images over time, as demonstrated in helical prominence studies where anti-symmetric Doppler patterns guide vector modeling. Observing solar prominences faces several challenges that complicate data interpretation. Limb darkening reduces contrast at the solar edge, making faint prominences harder to detect against the brighter disk interior. Projection effects along the obscure true three-dimensional structures, leading to ambiguities in flow directions and morphologies. Resolution limits, typically below 1 arcsecond for fine structures like threads (150–450 km wide), hinder detailed analysis, particularly for sub-arcsecond features in ground-based data affected by atmospheric turbulence.

Historical Development

Early Discoverals

The first detailed observations of solar prominences occurred during total solar eclipses in 1842 and 1843, when astronomers across reported bright, arc-like features projecting from the Sun's limb. During the eclipse of July 8, 1842, visible in , Italian astronomer Francesco de Vico at the Collegio Romano in , along with other observers such as in , Giuseppe Santini in , and English astronomer Francis Baily, described these reddish, flame-like or protuberance structures extending outward from the solar edge, visible even to the for brief moments after totality. These sightings marked the initial recognition of prominences as transient , distinct from the corona, though limited to eclipse conditions. Similar arc-shaped emissions were noted again during the December 21, 1843 eclipse, reinforcing their association with the Sun rather than the . Advancements in the mid-19th century revolutionized prominence studies by allowing observations beyond . In 1868, during the total of August 18 in , French astronomer and English astronomer Joseph Norman Lockyer independently developed a spectroscopic method to isolate the bright emission lines of prominences against the continuous solar disk , enabling their detection in full sunlight without waiting for an . This breakthrough identified prominences as dense, cool extensions of the rising into the hotter corona, composed primarily of and other elements emitting in the red Hα line at 656.3 nm. These observations also revealed an unknown yellow , later identified as , marking the first discovery of a new element on the Sun. Their reports, published concurrently, laid the groundwork for routine monitoring and confirmed prominences as integral to the solar atmosphere. By the early 20th century, systematic cataloging efforts provided deeper insights into prominence diversity. Between 1908 and 1912, , director of the , conducted extensive surveys using the newly invented spectroheliograph, documenting hundreds of prominences and classifying them based on morphological shapes—such as quiescent, active, and surge types—and dynamic behaviors like ascents, oscillations, and eruptions. These catalogs, comprising over 1,000 entries, highlighted patterns in prominence occurrence tied to solar activity cycles and facilitated the first quantitative analyses of their frequencies and evolutions. Hale's work shifted focus from sporadic glimpses to continuous, ground-based observation programs. Early interpretations of prominences often misconstrued them as permanent or semi-permanent upward extensions of the photospheric or chromospheric layers, without fully appreciating their embedded nature within the tenuous corona. Prior to spectroscopic confirmation, some astronomers even speculated they might pertain to a lunar atmosphere, though by the , evidence firmly established their solar origin; it was not until later coronal models in the 1930s that their suspended, cool-plasma state against the hotter surroundings was clarified, dispelling views of them as simple atmospheric protrusions.

Key Theoretical Advances

In the 1940s and 1950s, theoretical efforts shifted toward explaining how solar prominences could remain suspended in the hot corona, leading to the recognition that provide the necessary support against through the . Early explorations of electromagnetic forces in prominences laid the groundwork for this paradigm. A major milestone was the siphon model for active prominences, which described plasma flows along arched magnetic loops driven by differences between footpoints, facilitating mass supply and explaining the dynamic of prominences in active regions; this concept, developed in the late , built upon earlier magnetic support ideas from mid-century models. The Kippenhahn-Schlüter model of 1957 further advanced quiescent prominence by proposing a dense plasma sheet in a dipped magnetic arcade, where a current sheet generates an upward to balance , marking the first widely accepted static magnetohydrostatic configuration. During the 1960s and 1970s, models for quiescent prominences refined the dipped concept from the Kippenhahn-Schlüter framework, emphasizing inverse polarity configurations where field lines sag to trap cool plasma. Early magnetohydrodynamic (MHD) approaches, such as Pneuman's 1972 analysis of prominence magnetic structure, incorporated plasma condensation and curvature to explain stability and cavity formation around prominences. The Kuperus-Raadu model of 1974 extended this by depicting prominences as detached current filaments in a confining background field, supported by induced currents and suitable for both quiescent and intermediate types, highlighting the role of magnetic tension in long-term equilibrium. These analytical MHD efforts provided foundational insights into prominence threading and thermal isolation from the corona. In the 1980s and , theoretical focus turned to the role of prominences in processes, particularly their involvement in solar flares and eruptions. Prominences were integrated into reconnection models as sites where cancellation or shearing triggers energy release, with plasma dynamics facilitating reconnection in sheared arcades. Observations from the mission (1973–1974) provided critical data linking prominence activations to flares, revealing mass ejections and loop formations consistent with reconnection-driven models, such as those involving prominence material ejected during flare-associated events. This era saw high-impact contributions like the 1994 sheared arcade model by Antiochos et al., which demonstrated how photospheric shearing creates dipped fields primarily in inverse polarity configurations, supporting prominence threads while predisposing them to reconnection instabilities. Post-2010 advances have leveraged high-resolution data from the (SDO) to drive three-dimensional MHD simulations, uncovering complex inverse polarity configurations and multi-threaded structures in prominences. These models reveal how flux emergence and reconnection build twisted flux ropes that trap condensed plasma, with simulations reproducing observed barb-like features and providing quantitative constraints on magnetic helicity buildup. In the 2020s, insights from the have enhanced understanding of prominence-related plasma dynamics, including the in-situ measurement of helium-enriched ejecta from erupting prominences, confirming reconnection-driven acceleration and switch-on shocks in interplanetary space.

Classification

Types of Prominences

Solar prominences are categorized primarily according to their location relative to active regions, their typical lifetimes, and their associations with dynamic such as flares and coronal mass ejections (CMEs). This classification helps distinguish between stable, long-term structures and more transient, activity-linked features. Quiescent prominences represent the most stable class, persisting for weeks to months and often forming at high latitudes or as polar crown prominences encircling the Sun's poles. These structures are typically located far from s, in regions of weak magnetic activity, and show minimal association with eruptive events. In contrast, active region prominences are short-lived, lasting from hours to days, and form near sunspots within the strong of active regions. They are closely tied to solar activity, frequently preceding or accompanying flares and CMEs, with reported associations ranging from 3% to 90% depending on observational criteria. Surge prominences constitute a distinct, explosive subtype characterized by jet-like ejections of plasma from the solar surface, often originating in active regions. These events are brief, typically enduring only minutes, and involve high velocities that propel material outward before it often returns to the Sun. Beyond these primary categories, prominences exhibit various subtypes based on their apparent shapes and internal motions, including hedgerow prominences, which appear as linear arrangements of vertical threads with minimal plasma flows; fountain prominences, featuring loop-like structures with upward plasma flows transitioning to downward streams; and tornado prominences, which display rotating or vortex-like motions driven by helical , often linked to filament barbs and lasting from tens of minutes to hours. Classification relies on several key criteria to identify prominences as distinct from other solar features. Structures must extend to heights greater than 10,000 km above the to qualify, ensuring they are embedded in the corona rather than the . They are invariably associated with polarity inversion lines (neutral lines) where opposing magnetic polarities meet, providing the dipped field lines that support the cool plasma. Velocity profiles, derived from Doppler shifts, further aid differentiation, revealing internal flows typically ranging from 2 to 30 km/s, with counter-streaming or oscillatory patterns indicating the plasma's dynamic support.

Relation to Filaments

Solar filaments represent the same physical structures as solar prominences, consisting of cool, dense plasma suspended in the hot corona, but they are observed against the brighter background of the solar disk, appearing as dark, elongated features due to selective absorption of light, particularly in the (Hα) . This absorption occurs because the cooler plasma (approximately 5,000–10,000 K) in filaments blocks continuum radiation from the underlying more effectively than the surrounding . In contrast, when these structures are viewed at the solar limb, they manifest as prominences, appearing bright in emission against the dark sky due to the same Hα line but now in silhouette. The observational distinction arises primarily from line-of-sight geometry and projection effects. On the disk, filaments often exhibit an apparent disconnection or segmentation as overlying coronal structures or varying densities along the create illusory breaks in continuity. At the limb, the full three-dimensional structure of the prominence becomes visible, revealing loops or threads aligned with that support the plasma against gravity. This duality in appearance has led to the interchangeable use of terms, with "filament" reserved for on-disk views and "prominence" for limb views, emphasizing their unified nature as a single phenomenon. Historically, dark thread-like features on the solar disk were termed "filaments" as early as the late , based on visual observations during spectroheliography, while limb features were called "prominences" from studies dating back to the 1800s. The recognition that these were the same structures solidified in the early through spectroscopic analysis, with (1903) and Henri Deslandres (1910) demonstrating matching spectral signatures—such as strong Hα absorption/emission—of disk filaments and limb prominences, confirming their compositional identity. Further advances in post-1950s , including routine coronagraphic and filter-based observations, refined this understanding by revealing consistent plasma properties across both manifestations, dispelling earlier notions of separate origins. To distinguish true filaments—elevated, cool plasma threads—from pseudo-filaments, which may appear as dark linear features due to projection of chromospheric or other non-suspended structures, diagnostic employs Doppler shifts in lines like Hα or He I 10830 Å. These measurements reveal line-of-sight velocities (typically 5–20 km/s in downflows or counter-streaming flows) characteristic of dense, magnetized plasma in true filaments, whereas pseudo-filaments lack such organized, redshifted or blueshifted signatures indicative of mass support and dynamics.

Morphology

Structural Features

Solar prominences consist of cool, dense plasma organized into threads or barbs that are suspended within dips in the lines. These threads typically have temperatures between 7,000 K and 20,000 K and densities 10 to 100 times higher than the surrounding corona, allowing them to appear bright against the dark sky during limb observations or as dark filaments on the disk. Prominences form within filament channels, which are elongated regions in the and low corona aligned above photospheric polarity inversion lines—neutral lines where magnetic fields of opposite polarity meet. These channels exhibit aligned and serve as the base for the prominence structure, often spanning tens of thousands of kilometers along the neutral line. The internal includes a central spine acting as the primary axis, along which the main body of the prominence extends horizontally, and barbs that project perpendicularly from the spine like fine, thread-like extensions. Barbs, observed prominently in hedgerow-type prominences, have widths on the scale of 150–450 km and may connect downward toward the . Above the prominence lies an overlying cavity, a low-density that appears as a dark, elliptical envelope in (EUV) imaging, with brighter rims outlining its boundaries due to enhanced emission from the surrounding hotter plasma. Prominences exhibit a range of scales, with typical widths between 5,000 and 20,000 km, though their lengths can extend much farther along the filament channel. High-resolution imaging from instruments like the Interface Region Imaging Spectrograph (IRIS) has resolved fine-scale threads within the structure, down to 150–450 km, revealing dynamic variability and intricate plasma flows on small scales.

Magnetic Configurations

Solar prominences are supported by complex structures in the solar corona, primarily along polarity inversion lines (PILs) where photospheric of opposite polarity meets. The sheared arcade model, proposed by Antiochos et al., describes how prominences form in dipped lines within a sheared bipolar arcade, where photospheric shearing motions create concave-up field configurations capable of trapping cool plasma against gravity. In this model, ongoing cancellation at the PIL builds helical field components, allowing the prominence to reside in the magnetic dips while the overlying field arches provide stability. A key feature of many prominences is inverse polarity, where the direction in the prominence plasma is opposite to that of the surrounding coronal , often observed in magnetic concavities along the PIL. This configuration arises from twisted or tangled fields that reverse the polarity relative to the potential field extrapolation, enabling the plasma to accumulate in low-lying dips without direct support from vertical fields. Observational studies confirm that such inverse polarity is prevalent in quiescent prominences, with the field in the prominence threads typically horizontal and sheared. Prominences exhibit a consistent chirality pattern governed by Hale's law, with right-handed (dextral) fields dominating in the and left-handed (sinistral) fields in the , reflecting the hemispheric helicity rule observed in solar magnetic structures. This preference arises from the global dynamo and flux transport processes, leading to systematic twists in filament channels that align with the overall polarity. Alternative flux rope models describe quiescent prominences as twisted tubes embedded in the corona, where the helical fields provide both support and the potential for leading to eruptions. In these models, the axial field runs along the prominence, with azimuthal twists determining , and the rope's footpoints anchored near the PIL. Vector magnetograms from instruments like the Solar Dynamics Observatory's Helioseismic and Magnetic Imager (SDO/HMI) provide evidence for these structures, revealing horizontal magnetic fields of approximately 100–500 G along PILs that shear and support the overlying prominence configurations.

Formation

Quiescent Formation Processes

Quiescent solar prominences, also known as long-lived prominences, form through gradual processes that involve the reconfiguration of coronal magnetic fields and subsequent thermal evolution of plasma. These structures develop over extended periods, typically days, as cool, dense plasma accumulates in magnetic configurations that allow suspension against gravity. The primary mechanisms emphasize the interplay between and radiative processes, leading to the of chromospheric material into prominence threads. One key formation pathway is the reconnection-condensation model, in which small-scale events at the base of coronal arcades create dipped lines. These dips facilitate the drainage of cool plasma from the , where thermal instability promotes along the field lines. The thermal instability arises from an imbalance between heating and , causing coronal plasma to cool rapidly and increase in density, forming the cool core of the prominence. This process is supported by numerical simulations showing that reconnection triggers localized cooling, with condensed plasma flowing downward into the dips before stabilizing. An alternative or complementary mechanism is the levitation-condensation model, where magnetic buoyancy from emerging flux tubes lifts chromospheric plasma to coronal altitudes. Once elevated, the material undergoes and due to reduced heating at higher heights, accumulating in magnetic dips to form prominence threads. This model highlights the role of dynamic magnetic emergence in transporting low-temperature plasma against the hot corona. The timescales for quiescent prominence formation are characteristically slow, with buildup occurring over several days and times ranging from 10410^4 to 10510^5 seconds, governed by the plasma's thermal response to localized heating variations. For polar crown prominences, which are a subtype of quiescent structures, formation predominantly occurs at high solar latitudes (above 60°) during , where they are fed by cool plasma from unipolar magnetic regions at the boundaries of polar . These prominences often encircle the poles, forming a "crown" as part of the global reversal. Observational evidence for these processes includes detections of slow upflows in emission lines, such as those from the He II 304 Å and cooler transitions, observed by instruments like TRACE and SDO/AIA, indicating gradual mass accumulation during the initial stages of prominence development. These upflows, typically at speeds of a few km/s, align with the predicted drainage and levitation dynamics in dipped fields.

Active Region Formation

Solar prominences in s form rapidly within magnetically complex environments characterized by strong, sheared and frequent flux emergence. Unlike slower quiescent processes, these formations occur amid dynamic interactions between emerging and preexisting coronal structures, often leading to the buildup of cool, dense plasma threads suspended in dipped . This rapid development is closely tied to the evolution of sunspots and filaments within the active region, contributing to heightened solar activity. One key mechanism involves siphon flows, where asymmetric heating at the footpoints of arched magnetic loops drives plasma circulation. Chromospheric heating on one leg of the loop causes and upflow, while the cooler plasma drains along the other leg toward the loop apex, accumulating dense material in magnetic dips. This process can sustain prominence growth even after initial heating subsides, with prominences exhibiting longer thread lengths due to stronger compared to quiescent cases. Simulations show that once a seed condensation forms at the apex, siphon flows from both footpoints inject additional mass, enhancing the structure's stability. Another prominent mechanism is reconnection injection, triggered by emerging that interacts with overlying fields to create sheared configurations. As bipoles emerge through the , magnetic at the interface forms current sheets and ejects multithermal jets, which cool and drain into filament channels. Repetitive reconnection events progressively inject cool chromospheric plasma into these channels, building the prominence body over short intervals. This process is particularly efficient in active regions, where the emerging flux introduces helicity and supports inverse polarity fields conducive to prominence support. Recent unified models, incorporating self-consistent heating, combine evaporation-condensation (for upper chromospheric heating leading to plasma evaporation and subsequent cooling) with direct injection (for lower chromospheric heating driving pressure-induced plasma rise), providing a comprehensive framework for these dynamics. These formations typically occur on timescales of hours to a few days, directly linked to emergence rates of approximately 101810^{18}–101910^{19} per hour in s. Initial flux emergence can initiate plasma within hours, with full prominence development following as reconnection and flows accumulate mass over 1–2 days in observed cases. Such rapid buildup contrasts with longer quiescent timescales and underscores the role of flux emergence in driving dynamics. Prominences in active regions are frequently associated with complex sunspot configurations, particularly near delta spots where umbrae of opposite polarity overlap within a shared penumbra. Here, filament barbs—short, perpendicular extensions from the main prominence axis—often align with underlying penumbral filaments, reflecting the twisted that anchor the structure. This alignment facilitates mass support and is evident in regions with high shear, enhancing the prominence's integration into the active region's magnetic arcade. Observations from the (SDO) in 2011, such as those of NOAA 11158, captured prominence formation preceding X-class flares, including the cycle's first X2.2 event on February 15. These events highlighted how emerging flux and reconnection in the region led to rapid plasma accumulation hours to days before flare onset, linking prominence buildup to subsequent eruptive activity.

Dynamics

Stability Mechanisms

Solar prominences maintain their structure through a delicate balance of forces, primarily supported against by magnetic tension in dipped lines. The , arising from the interaction between plasma currents and the , provides the upward component necessary to counteract gravitational pull. This equilibrium is described by the magnetohydrostatic equation J×B=∇p−ρg\mathbf{J} \times \mathbf{B} = \nabla p - \rho \mathbf{g}, where the vertical component of the balances the plasma pressure gradient and gravitational force, with J\mathbf{J} as the , B\mathbf{B} the , pp the pressure, ρ\rho the density, and g\mathbf{g} . In dipped field configurations, such as those formed by sheared arcades, the magnetic tension acts like a hammock, suspending the cool, dense plasma in the corona. The Kippenhahn-Schlüter (KS) model provides an analytical framework for this equilibrium, representing a prominence as an infinite slab embedded in a vertical magnetic sheet within a quadrupolar field. In this model, a current sheet supports the plasma, with the upward exactly balancing , leading to a stable configuration for uniform and . This model assumes isothermal conditions and neglects horizontal variations, offering a foundational understanding of how magnetic tension sustains prominences over extended periods. Thermal stability in prominences relies on a balance between radiative losses and heating mechanisms, preventing runaway cooling or . , dominated by recombination and line emission in the cool plasma (temperatures around 7500–9000 K), is offset by coronal heating sources—potentially including along field lines and wave dissipation—maintaining the prominence's temperature and density. This energy equilibrium is crucial for , as imbalances can trigger thermal instabilities, though steady states allow prominences to persist. Prominences exhibit potential instabilities, such as the kink mode in twisted flux ropes, with growth timescales on the order of hours based on Alfvén transit times. However, these are often suppressed by line-tying effects, where magnetic footpoints anchored in the dense resist perturbations, stabilizing the structure against helical deformations. Observationally, stability manifests as quasi-static evolution over days to weeks, with prominences showing minimal large-scale motion and only minor oscillations (amplitudes of a few km/s, periods from minutes to hours) driven by flows or waves. These indicators, detected via Doppler shifts and high-resolution imaging, confirm the robustness of the magnetic and thermal equilibria.

Eruption and Evolution

Solar prominence eruptions are initiated by mechanisms that disrupt the delicate balance of magnetic forces supporting the structure. A primary trigger is the critical Lorentz imbalance, where the outward Lorentz force from internal magnetic pressure and plasma gradients overcomes the downward tension provided by dipped magnetic field lines and overlying coronal fields, leading to upward motion. Tether-cutting reconnection, occurring in the sheared fields above the prominence, severs the anchoring flux and releases stored magnetic energy, propelling the prominence outward. Additionally, flux cancellation at the photospheric footpoints reduces the magnetic tension holding the prominence in place, gradually destabilizing the configuration and contributing to eruption onset. Instabilities play a crucial role in driving the ejection process once initial perturbations occur. The kink instability develops from helical perturbations in highly twisted flux ropes, where the magnetic twist exceeds a critical threshold of approximately 2.5π radians (1.25 turns), causing the structure to writhe and accelerate away from the solar surface. Similarly, the torus instability arises when the overlying decays too rapidly with height, amplifying outward expansion through toroidal forces and leading to rapid ejection; this process is often intertwined with kink modes in twisted configurations. These instabilities transform gradual disruptions into dynamic ejections, with helical deformations observable in high-resolution imaging. The evolution of erupting prominences unfolds in distinct phases. An initial slow-rise phase sees the structure ascending at velocities of 10–50 km/s over tens of minutes to hours, driven by partial loss of equilibrium. This transitions to a rapid phase, where speeds reach up to 1000 km/s, propelled by reconnection and instability growth, culminating in either partial eruptions—where fragments remain near the Sun—or full ejections into the . Studies indicate that 50–90% of these eruptions are associated with coronal mass ejections (CMEs), with one analysis finding approximately 70%, which carry prominence material into space, and many coincide with solar flares due to the sudden release of . Following an eruption, prominences can undergo reformation through secondary reconnection events. In 3D magnetohydrodynamic (MHD) simulations, reconnected fields from the interaction of the ejected flux rope with surrounding arcade structures form new dipped fields, allowing cool plasma to condense and rebuild a prominence-like feature at lower heights. This process highlights the dynamic recycling of magnetic structures in the solar atmosphere, often observed in post-eruptive arcades.

Significance

Role in Solar Activity

Solar prominences exhibit a pronounced dependence on the solar activity cycle, with their occurrence rates peaking during , mirroring the behavior of sunspots and other magnetic phenomena. Observations spanning multiple cycles reveal a characteristic time-latitude distribution akin to the , where high-latitude polar prominences emerge around 40°–50° during cycle minima and display poleward migration that accelerates near 70° latitude, reaching rates of up to 19.5 m/s in the high polar zone during the rising phase. This pattern reflects the underlying processes driving reversal at the poles. The formation and persistence of prominences signal significant magnetic complexity in the solar atmosphere, often indicating the accumulation of opposite-polarity along polarity inversion lines that precede or accompany the of active regions. In active regions, prominences typically develop within sheared magnetic configurations where flux buildup creates dipped field lines capable of supporting cool plasma against , serving as early markers of heightened magnetic stress. Prominences function as key reservoirs of relatively cool, dense plasma embedded in twisted , providing material that can fuel solar flares and coronal mass ejections (CMEs) upon destabilization and eruption. Their role extends to transport, as evolving prominences facilitate reconnection and cancellation processes that redistribute flux across latitudes, contributing to the global evolution of the solar magnetic field. Comprehensive catalogs from and SDO missions document roughly 10510^5 prominences over an 11-year cycle, with 20–30% undergoing eruptions that link to broader activity. Recent analyses from the 2020s, leveraging SDO data through , highlight how prominence disappearance rates—often via activation or partial eruption—correlate with cycle phase, increasing during the descending phase as polar field reversal nears completion and magnetic reconfiguration intensifies. These studies underscore prominences' integral role in cycle progression, with disappearance events peaking in tandem with declining numbers. As of November 2025, remains in its maximum phase, with continued prominence activity monitored by missions like , revealing enhanced eruption rates linked to complex active regions.

Impacts on Space Weather

Erupting solar prominences frequently serve as the dense plasma core of coronal mass ejections (CMEs), driving significant space weather effects when directed toward Earth. Studies indicate that prominence eruptions are associated with approximately 70% of all observed CMEs, with the association rate reaching up to 94% for Earth-directed halo CMEs, which are particularly geoeffective. These CMEs typically eject 10^{15} to 10^{16} grams of magnetized plasma into the heliosphere, with the prominence material contributing substantially to the core density and structure. Upon reaching after 1 to 4 days of travel, these CMEs interact with the , compressing it and inducing s that enhance auroral activity visible at low latitudes and generate (GICs). These GICs can overload power grids, as exemplified by the severe of November 8–10, 1991, triggered by a massive filament (prominence) eruption spanning 25° in , which produced a CME with speeds exceeding 1800 km/s and led to widespread ionospheric scintillation and vulnerabilities. Such events underscore the role of prominence-driven CMEs in disrupting electrical infrastructure, with historical analyses showing GIC amplitudes up to several amperes per kilometer in vulnerable regions. Prominence eruptions also accelerate (SEPs) to near-relativistic speeds, increasing radiation hazards for satellites and high-altitude aircraft while causing ionospheric disturbances that degrade GPS accuracy through signal scintillation and phase delays. SEP fluxes from these events can exceed 10^2 particles cm^{-2} s^{-1} sr^{-1} above 10 MeV, necessitating operational pauses for unshielded and enhanced monitoring for systems. Monitoring prominence instabilities and early CME signatures with instruments like the Large Angle and Spectrometric Coronagraph (LASCO) on the () enables forecasting of Earth-impacting events, providing advance warnings of 1–4 days based on CME propagation speeds from 300 to 2000 km/s. This observational capability has improved alerts, allowing mitigation measures such as reconfiguration. In 2024, during , multiple prominence eruptions contributed to intense CME activity, including those from AR 3664 in May, which elevated radiation levels to S1 (minor) with spikes to S2 (moderate) and prompted precautions for the crew to limit exposure to SEPs. These events highlighted the ongoing need for adaptive shielding protocols amid heightened solar activity.

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