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Chromosphere
Chromosphere
Chromosphere
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When observed in the Hα spectral line, the chromosphere appears deep red.

A chromosphere ("sphere of color", from the Ancient Greek words χρῶμα (khrôma) 'color' and σφαῖρα (sphaîra) 'sphere') is the second layer of a star's atmosphere, located above the photosphere and below the solar transition region and corona. The term usually refers to the Sun's chromosphere, but not exclusively, since it also refers to the corresponding layer of a stellar atmosphere. The name was suggested by the English astronomer Norman Lockyer after conducting systematic solar observations in order to distinguish the layer from the white-light emitting photosphere.[1][2]

In the Sun's atmosphere, the chromosphere is roughly 3,000 to 5,000 kilometers (1,900 to 3,100 miles) in height, or slightly more than 1% of the Sun's radius at maximum thickness. It possesses a homogeneous layer at the boundary with the photosphere. Narrow jets of plasma, called spicules, rise from this homogeneous region and through the chromosphere, extending up to 10,000 km (6,200 mi) into the corona above.

The chromosphere has a characteristic red color due to electromagnetic emissions in the Hα spectral line. Information about the chromosphere is primarily obtained by analysis of its emitted electromagnetic radiation.[3] The chromosphere is also visible in the light emitted by ionized calcium, Ca II, in the violet part of the solar spectrum at a wavelength of 393.4 nanometers (the Calcium K-line).[4]

Chromospheres have also been observed on stars other than the Sun.[5] On large stars, chromospheres sometimes make up a significant proportion of the entire star. For example, the chromosphere of supergiant star Antares has been found to be about 2.5 times larger in thickness than the star's radius.[6]

Physical properties

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The red color of the chromosphere could be seen during the solar eclipse of August 11, 1999.

The density of the Sun's chromosphere decreases exponentially with distance from the center of the Sun by a factor of roughly 10 million, from about 2×10−4 kg/m3 at the chromosphere's inner boundary to under 1.6×10−11 kg/m3 at the outer boundary.[7] The temperature initially decreases from the inner boundary at about 6000 K[8] to a minimum of approximately 3800 K,[9] but then increasing to upwards of 35,000 K[8] at the outer boundary with the transition layer of the corona (see Stellar corona § Coronal heating problem).

The density of the chromosphere is 10−4 times that of the underlying photosphere and 10−8 times that of the Earth's atmosphere at sea level. This makes the chromosphere normally invisible and it can be seen only during a total eclipse, where its reddish colour is revealed. The colour hues are anywhere between pink and red.[10] Without special equipment, the chromosphere cannot normally be seen due to the overwhelming brightness of the photosphere.

The chromosphere's spectrum is dominated by emission lines when observed at the solar limb.[11][12] In particular, one of its strongest lines is the Hα at a wavelength of 656.3 nm; this line is emitted by a hydrogen atom whenever its electron makes a transition from the n=3 to the n=2 energy level. A wavelength of 656.3 nm is in the red part of the spectrum, which causes the chromosphere to have a characteristic reddish colour.

Phenomena

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Further information: Solar phenomena
High-resolution observations of the solar chromosphere show hair-like spicules, here shown in a false colored image made in borderline ultraviolet radiation of calcium K-line.

Many different phenomena can be observed in chromospheres.

Plage

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Main article: Solar plage

A plage is a particularly bright region within stellar chromospheres, which are often associated with magnetic activity.[13]

Spicules

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Main article: Solar spicule

The most commonly identified feature in the solar chromosphere are spicules. Spicules rise to the top of the chromosphere and then sink back down again over the course of about 10 minutes.[14]

Oscillations

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Since the first observations with the instrument SUMER on board SOHO, periodic oscillations in the solar chromosphere have been found with a frequency from 3 mHz to 10 mHz, corresponding to a characteristic periodic time of three minutes.[15] Oscillations of the radial component of the plasma velocity are typical of the high chromosphere. The photospheric granulation pattern usually has no oscillations above 20 mHz; however, higher frequency waves (100 mHz, or a 10 s period) were detected in the solar atmosphere (at temperatures typical of the transition region and corona) by TRACE.[16]

Loops

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See also: Coronal loop

Plasma loops can be seen at the border of the solar disk in the chromosphere. They are different from solar prominences because they are concentric arches with maximum temperature of the order 0.1 MK (too low to be considered coronal features). These cool-temperature loops show an intense variability: they appear and disappear in some UV lines in a time less than an hour, or they rapidly expand in 10–20 minutes. Foukal[17] studied these cool loops in detail from the observations taken with the EUV spectrometer on Skylab in 1976. When the plasma temperature of these loops becomes coronal (above 1 MK), these features appear more stable and evolve over longer times.

Network

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Images taken in typical chromospheric lines show the presence of brighter cells, usually referred to as the network, while the surrounding darker regions are named internetwork. They look similar to granules commonly observed on the photosphere due to the heat convection.

On other stars

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Chromospheres are present on almost all luminous stars other than white dwarfs. They are most prominent and magnetically active on lower-main sequence stars, on brown dwarfs of F and later spectral types, and on giant and subgiant stars.[13]

A spectroscopic measure of chromospheric activity on other stars is the Mount Wilson S-index.[18][19]

See also

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References

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Revisions and contributorsEdit on WikipediaRead on Wikipedia

Chromosphere

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The chromosphere is the thin, irregular second layer of the Sun's atmosphere, positioned immediately above the and extending to the base of the transition region before the corona. Approximately 2,000 kilometers thick, it features a plasma density about 10,000 times lower than that of the photosphere, with temperatures rising sharply from roughly 6,000 at its base to around 20,000 at the top. This layer derives its name from words for "color" and "sphere," reflecting its vivid reddish appearance during total solar eclipses, caused by strong emissions in the at 656.3 nanometers. It is also prominently observed in the ionized calcium (Ca II) H and K lines at 393.4 nanometers and 396.8 nanometers, which reveal fine-scale structures like the chromospheric network—a web-like pattern of supergranulation boundaries. The chromosphere is a highly dynamic dominated by the Sun's , where plasma motions drive rapid changes over minutes, including the formation of spicules—short-lived, jet-like eruptions—and the chromospheric network of plage regions associated with magnetic activity. It hosts cool, dense plasma structures such as filaments and prominences, which appear as dark features against the brighter disk in images and can extend thousands of kilometers while remaining suspended by . Solar flares and prominence eruptions often originate here, contributing to events like coronal mass ejections that impact Earth's .

Definition and Location

Position in the Solar Atmosphere

The solar atmosphere is structured in layers extending outward from the Sun's visible surface, beginning with the as the innermost layer, followed by the chromosphere, a thin transition region, and the expansive corona. The chromosphere occupies the position immediately above the , serving as the intermediate layer between this dense, optically thick base and the hotter, more tenuous outer atmosphere. This positioning makes it a critical boundary zone in the solar atmosphere, bridging the cooler surface regions with the plasma-dominated upper layers. As a partially ionized plasma, the chromosphere represents a transitional domain where neutral atoms coexist with ions and electrons, influencing energy transport and magnetic interactions across the atmosphere. It is in this layer that the temperature profile reaches a minimum near its base, after which it begins to increase outward, setting the stage for the dramatic heating observed in the overlying corona. This temperature inversion marks the chromosphere's role in the overall thermal structure of the solar atmosphere, providing essential context for understanding plasma dynamics and radiative processes in subsequent layers. The chromosphere is approximately 2,000 km thick, extending from about 500 km to 2,500 km above the in its quiescent state, though dynamic extensions through structures like spicules can reach up to 10,000 km; further details on such structures are addressed elsewhere. These values are based on semi-empirical models like the VAL C and may vary slightly depending on the quiet-Sun region considered.

Boundaries and Thickness

The chromosphere's lower boundary is defined at the temperature minimum, located approximately 500 km above the , where temperatures reach a minimum of about 4500 K. This marks the transition from the cooler upper to the heating chromosphere. The upper boundary lies at the base of the transition region, a narrow zone roughly 2500 km above the , where temperatures abruptly rise from around 20,000 K to exceed 10^5 K over a mere 100 km thickness. The chromosphere's average thickness is approximately 2000 km, though this dimension varies spatially and temporally due to influences from solar activity, such as configurations that can extend its effective reach through dynamic processes. Across these boundaries, density decreases by several orders of magnitude, from photospheric levels to much lower coronal values. Defining precise boundaries remains challenging owing to the chromosphere's highly dynamic nature, with rapid temporal and spatial changes that blur static height demarcations, compounded by observational limitations in resolving sub-arcsecond structures from ground- or space-based telescopes. Historical measurements originated from total solar eclipses, where the flash spectrum first revealed the chromosphere's extent; in 1868, analyzed eclipse observations to name the layer and estimate its scale based on emission line altitudes. Modern observations, leveraging instruments like those on the , provide refined altitude profiles through multi-wavelength imaging, surpassing early eclipse-based geometric estimates.

Physical Properties

Temperature and Density Profiles

The temperature profile of the solar chromosphere is characterized by a non-monotonic , beginning at approximately 6,000 immediately above the . This decreases to a minimum of about 4,200 at a height of roughly 500 km according to semi-empirical models like VAL C, though recent simulations indicate values as low as 2,800 , after which the temperature rises abruptly to around 20,000 in the upper layers adjacent to the corona. These values are established through semi-empirical modeling constrained by EUV observations of the quiet Sun, while contemporary 3D simulations reveal spatiotemporal fluctuations driven by acoustic and magnetohydrodynamic waves, yielding local minima below 3,000 and more nuanced insights into energy deposition as of 2024. The mass density profile exhibits an exponential decline from approximately 2×10−42 \times 10^{-4} kg/m³ at the chromospheric base to about 10−1210^{-12} kg/m³ at the top, spanning a range of 2,000–2,500 km. This steep falloff arises from the condition of , expressed as dPdh=−ρg,\frac{dP}{dh} = -\rho g, where PP denotes gas , ρ\rho is mass , g≈274g \approx 274 m/s² is the effective near the solar surface, and hh is above the . In the approximation, relates to and via P∝ρTP \propto \rho T, thereby coupling the thermal and density structures. These profiles critically influence and energy balance within the chromosphere. Higher densities in the lower layers promote collisional excitation and efficient cooling through bound-bound and bound-free transitions, whereas the dilute upper atmosphere approaches optically thin conditions, enhancing emission in lines like Hα and Ca II. Seminal 1D models such as VAL C capture the mean structure, but contemporary 3D simulations demonstrate spatiotemporal fluctuations driven by acoustic and magnetohydrodynamic waves, yielding more nuanced insights into local energy deposition and dissipation.

Composition and Ionization

The of the chromospheric plasma closely mirrors that of the underlying , with accounting for approximately 92% of atoms by number, about 8%, and metals comprising less than 1% (primarily oxygen, carbon, , iron, and others in trace amounts). These abundances reflect the solar nebula's primordial mix, largely unaltered in the lower and middle chromosphere, though subtle enhancements in low first-ionization-potential (low-FIP) elements like iron and magnesium can occur in active regions due to plasma fractionation processes, with IRIS observations confirming factors up to 4× in plages as of 2023. At the base of the chromosphere, adjacent to the , the plasma remains predominantly neutral, dominated by atomic (H I) and (He I), with fractions for below 0.1 under temperatures below ~10,000 K. As height and temperature increase toward the upper chromosphere—reaching up to around 20,000 K—the state evolves significantly, with partial (0.1–0.9) in the middle layers at T up to ~17,000 K, transitioning to fully ionized protons (H II) at temperatures exceeding 23,000 K and ionizing to He II. Metals, present in minute abundances, are chiefly singly ionized throughout the layer (e.g., Ca II, Mg II, Fe II), contributing electrons that enhance conductivity and influence , while higher stages like Fe III appear only in hotter, dynamic regions. Recent studies highlight non-equilibrium effects in the partially ionized plasma, affecting heating and dynamics as of 2024. The balance plays a in determining the chromosphere's opacity, which governs how radiation escapes the plasma, and in producing diagnostic emission lines. Neutral , for example, absorbs and re-emits photons in the , with the prominent Hα line (at 656.3 nm) arising from the transition between the n=3 and n=2 principal quantum levels in H I, providing insights into mass motions and heating. Singly ionized metals contribute strong resonance lines, such as the Ca II H and K lines or Mg II h and k lines, whose formation depends on the local and temperature, enhancing the layer's visibility in and optical spectra. These ionization profiles vary sharply with the steep temperature gradients in the chromosphere, from near-neutral conditions at the base to highly states aloft, with recombination times shortening from hundreds of seconds in cooler zones to under 10 seconds in hotter ones. Recent spectroscopic data from missions like Hinode's Extreme-ultraviolet Imaging Spectrometer (EIS) and the Interface Region Imaging Spectrograph (IRIS) have updated abundance determinations, confirming photospheric-like compositions in quiet-Sun regions while highlighting localized low-FIP enhancements (up to factors of 2-4 for elements like Fe) tied to chromospheric heating and wave activity.

Structure and Dynamics

Magnetic Field Role

The solar chromosphere is permeated by ubiquitous that play a central role in its structuring and dynamics. In quiet-Sun regions, these fields exhibit strengths of approximately 10–100 G at the base of the chromosphere, increasing to around 1,000 G in network regions where magnetic concentrations are prominent. These field strengths reflect the transition from photospheric origins to chromospheric amplification, influencing plasma behavior across the layer. Magnetic flux tubes serve as the primary structural elements within the chromosphere, confining and guiding plasma flows while facilitating energy transport from lower atmospheric layers. These tubes, often rooted in the , channel material and disturbances upward, maintaining the dynamic equilibrium of the plasma against gravitational and thermal forces. The influence of these fields is quantified by the , given by Pmag=B22μ0,P_{\text{mag}} = \frac{B^2}{2 \mu_0}, where BB is the magnetic field strength and μ0\mu_0 is the permeability of free space. In the upper chromosphere, this magnetic pressure becomes comparable to the gas pressure, marking a regime where plasma β≈1\beta \approx 1 and magnetic forces significantly shape the plasma dynamics. Magnetic fields drive chromospheric heating primarily through mechanisms such as magnetic reconnection and magnetohydrodynamic (MHD) waves, which dissipate energy into thermal form. Reconnection events release stored magnetic energy, accelerating particles and heating plasma locally, while Alfvén and magnetoacoustic waves propagate along flux tubes, undergoing mode conversion and damping to contribute to the overall energy budget. Three-dimensional MHD simulations have been instrumental in elucidating these processes, demonstrating how wave dissipation and reconnection in realistic flux tube geometries account for observed heating rates in the quiet chromosphere.

Layered Substructure

The solar chromosphere exhibits a layered substructure distinguished by variations in physical conditions, density, and dominant dynamic processes, extending approximately from 0 to 2,500 km above the . This division into lower, middle, and upper regions arises from empirical modeling and spectroscopic observations that reveal transitions in atmospheric behavior, with the lower chromosphere influenced primarily by , the middle by wave-driven dynamics, and the upper by emerging magnetic influences leading to the corona. The lower chromosphere, spanning roughly 0 to 1,000 km in height, is characterized by its close coupling to the underlying convective motions from the , where patterns extend upward and drive propagation. In this region, the plasma remains relatively dense and optically thick, with acoustic waves generated by photospheric propagating without significant steepening, maintaining oscillatory behavior that contributes to local heating. Empirical models indicate that this layer forms a bridge between the cooler and higher regions, with minimal dominance allowing convective influences to prevail. In the middle chromosphere, approximately 1,000 to 2,000 km above the , acoustic waves from below steepen into shock waves due to decreasing , leading to enhanced heating and such as fibril-like configurations. This layer features increased inhomogeneity, with shock formation periods on the order of 10–200 seconds, resulting in intermittent energy dissipation that shapes the overall thermal structure. Observations and simulations highlight the role of these shocks in maintaining the plateau observed here, distinct from the smoother propagation in lower heights. The upper chromosphere, extending above 2,000 km up to the transition region around 2,500 km, marks a shift toward dominance, where open field lines facilitate the acceleration of plasma outflows and the onset of coronal conditions. Dynamic processes in this layer include intensified shock interactions and , pre-conditioning the atmosphere for the million-degree corona, with reduced convective input allowing magnetohydrodynamic effects to govern structure. This region's characteristics are evident in emissions tracing the thinning plasma. These sublayers are primarily delineated through semi-empirical models like the Vernazza-Avrett-Loeser (VAL) series, particularly the VAL-C model for the quiet Sun, which integrates brightness observations from to construct height-dependent profiles of temperature, density, and ionization, revealing the progressive changes across the chromosphere without assuming a uniform structure. Updates to such models continue to refine these divisions based on modern spectroscopic data, emphasizing the chromosphere's role as a dynamically evolving interface.

Observable Phenomena

Spicules and Fibrils

Spicules are dynamic, needle-like jets of plasma that extend from the solar photosphere into the chromosphere, typically measuring 300–500 km in diameter and reaching lengths of 5,000–15,000 km. These structures have lifetimes of approximately 5–15 minutes and exhibit upward velocities ranging from 10–30 km/s, often following parabolic trajectories indicative of ballistic motion. Spicules are guided along lines, channeling plasma motions within the chromosphere's magnetized environment. Two primary types of spicules have been identified based on their formation mechanisms and dynamics. Type I spicules are driven by shock waves propagating from the , resulting in slower, more prolonged ejections with velocities of 15–40 km/s and lifetimes of 150–400 seconds. In contrast, Type II spicules arise from events, forming rapidly with higher velocities of 30–150 km/s and shorter lifetimes of 10–150 seconds, often fading without a clear descent phase as the plasma heats and disperses. Observations from the Interface Region Imaging Spectrograph (IRIS) have revealed that Type II spicules undergo significant heating to transition region temperatures during their ascent, distinguishing them further from the cooler Type I variety. Fibrils represent horizontal, elongated extensions of spicular plasma, appearing as dark, thread-like features in chromospheric lines such as Hα, and typically aligning along magnetic neutral lines where opposing field polarities meet. These structures, with widths around 700 km and lengths up to 14,000 km, exhibit swaying motions and serve as tracers of horizontal magnetic fields in the lower chromosphere, often connecting regions of enhanced magnetic activity. Spicules, particularly Type II, contribute substantially to chromospheric heating and the supply of mass and energy to the corona through their kinetic and thermal fluxes. IRIS observations indicate that these spicules transport plasma at rates sufficient to contribute to the corona's mass and energy supply through their kinetic and thermal effects. This process underscores spicules' role in bridging the chromosphere and corona, facilitating the transfer of photospheric energy upward.

Plages and Network

Plages are bright regions in the solar chromosphere, appearing as enhanced emission patches in the Hα line at 6563 and the Ca II K line at 3934 , typically spanning 10–20 arcseconds in angular size. These features are closely associated with concentrations of , often manifesting as bipolar magnetic regions where opposite polarities emerge and interact, leading to heightened chromospheric activity. The increased of plages arises from elevated temperatures, reaching up to 10,000 K compared to the quieter chromosphere's 4,000–6,000 K, which intensifies line emission through collisional excitation and . The chromospheric network forms a pervasive cellular across the solar disk, delineated by elongated lanes of enhanced emission that trace the boundaries of supergranular cells. These lanes consist of intergranular concentrations, with typical spacing between network elements around 30,000 km, reflecting the scale of underlying supergranulation. The network's persistence is maintained by the horizontal convective flows of supergranulation, which advect photospheric toward cell boundaries, concentrating it into stable, elongated structures that extend upward into the chromosphere. Both plages and the chromospheric network exhibit pronounced variations over the 11-year , with their coverage and intensity peaking during periods of maximum solar activity. In contrast, during , these features diminish, revealing a more subdued quiet-Sun network sustained primarily by residual diffusion and weak .

Oscillations and Waves

The solar chromosphere exhibits prominent oscillations driven by originating from the underlying . These include the well-known 3-minute oscillations, which correspond to p-modes with periods ranging from approximately 180 to 300 seconds, propagating upward as acoustic or magneto-acoustic waves into the chromospheric layers. These waves are particularly dominant in the internetwork regions of the quiet Sun, where they manifest as intensity and velocity fluctuations observed in chromospheric lines such as Hα and Ca II K. The propagation of these p-modes is influenced by the decreasing density with height, leading to wave steepening and potential shock formation that contributes to atmospheric structuring. In contrast, 5-minute oscillations, associated with global p-modes of around 300 seconds period, are more pronounced at the boundaries of the chromospheric network, where channel their energy leakage into the upper atmosphere. These longer-period waves, evanescent in the non-magnetic internetwork, propagate along inclined tubes in regions, enhancing oscillatory power and influencing the dynamics at supergranular boundaries. Recent high-resolution observations from the (DKIST) have revealed detailed structures of chromospheric swirls in H-alpha, enhancing understanding of their role in exciting Alfvén waves. This spatial differentiation highlights how magnetic topology modulates the transmission of convective oscillations from the to the chromosphere. Beyond acoustic modes, magnetohydrodynamic (MHD) waves, particularly Alfvén waves, play a crucial role in transporting non-thermal energy through the chromosphere. Alfvén waves, incompressible transverse perturbations guided by magnetic tension, propagate at the Alfvén speed given by vA=Bμ0ρ,v_A = \frac{B}{\sqrt{\mu_0 \rho}},
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