The chromosphere is a roughly 2,000-kilometer-thick layer of the Sun’s atmosphere sandwiched between the visible surface (the photosphere) and the searingly hot corona above. Its name comes from the Greek word for “color,” because during total solar eclipses it briefly flashes into view as a thin reddish arc around the Moon’s silhouette. What makes this region so puzzling, and so central to modern solar physics, is that its temperature climbs as you move away from the Sun’s surface, rising from a minimum near 3,800 K to tens of thousands of kelvin at its upper boundary, defying the intuition that things should cool down the farther they get from a heat source.
Where the Chromosphere Sits and What It Looks Like
The photosphere, which is the layer you see when you look at the Sun through a safe filter, has a temperature of about 5,800 K. Moving outward, the temperature dips before it starts climbing again. Atmospheric models place this minimum at around 3,800 K, corresponding to a gas pressure of roughly 80 dyne per square centimeter. That minimum-temperature layer marks the base of the chromosphere. Above it, the temperature rises steeply through the chromosphere, eventually reaching roughly 10,000 to 20,000 K in its upper reaches.
Then comes one of the sharpest transitions in all of astrophysics. A narrow zone called the transition region, less than about 200 kilometers thick in the real Sun, separates the chromosphere from the corona. Within that sliver, the temperature rockets from around 12,000 K to about 600,000 K, while the density drops by two orders of magnitude. Simulations reproduce this jump, though they tend to stretch the transition region somewhat wider than what’s actually observed.
The Temperature Puzzle
In any normal situation, a gas should cool as it moves farther from a heat source. The chromosphere breaks that rule. From the temperature minimum upward, the gas gets hotter, not cooler, and the corona above is hotter still, reaching millions of kelvin. This “coronal heating problem” has been one of the biggest open questions in astrophysics for decades, but the chromosphere’s own heating is just as important and arguably harder to explain, because the chromosphere is far denser than the corona and radiates energy away much more efficiently. Whatever mechanism heats this layer has to pump in a lot of energy just to keep up with those losses.
Recent simulations show that a temperature inversion like the Sun’s can arise naturally in a gravitationally bound plasma when certain conditions are met, with the transition region forming as a sharp boundary between the cooler chromospheric gas below and the hot, tenuous corona above. The simulated temperature and density profiles are strikingly similar to the real Sun’s atmosphere, lending support to models that treat the heating as an emergent property of the plasma rather than something requiring a single exotic mechanism.
Acoustic Waves and Shock Heating
One of the leading candidates for chromospheric heating is acoustic-wave energy generated in the convective zone below the photosphere. Turbulent convective motions produce sound waves that travel upward. Because the density drops with altitude, these waves steepen into shocks as they rise, and the energy they dissipate heats the surrounding gas. In simulations of the quiet Sun (areas away from active regions), the friction at the shock front produces viscous dissipation that far exceeds the heating from electrical resistance, with the viscous component sometimes reaching more than six times the ohmic component and often peaking above 10 kilowatts per square meter.
That said, observations suggest the story isn’t quite so simple. Comparisons of the acoustic energy flux deposited in magnetic areas and plage regions against the radiative cooling of those same regions show that the acoustic contribution, while significant, falls short of balancing the losses. Part of the gap may be observational: we can’t yet measure waves at all the relevant frequencies, and the transfer of wave energy into the chromosphere isn’t perfectly understood. But the shortfall hints that additional heating mechanisms, possibly involving magnetic-field interactions, are at work in magnetically active parts of the chromosphere.
Spicules and the Chromosphere’s Restless Surface
If you could zoom in on the chromosphere with a powerful enough telescope, you would not see a smooth, featureless shell. Instead, you’d see a forest of thin, jet-like structures called spicules shooting upward from the surface at any given moment. These features are among the most visually striking aspects of solar physics, and they may play a role in channeling energy and mass into the corona.
Spicules come in two broad categories. Type I spicules move relatively slowly, at about 10 to 40 kilometers per second, and last anywhere from 3 to 10 minutes before falling back. Type II spicules are faster and more fleeting, with speeds ranging from roughly 80 to 300 kilometers per second and lifetimes of only 1 to 3 minutes. Type II spicules are especially interesting because they appear to inject some material into the transition region and low corona. Most of the spicule material, though, is only weakly heated; it shows up prominently in extreme-ultraviolet light as it falls back toward the surface rather than staying hot enough to join the corona permanently.
Recent high-resolution observations have added more complexity. Researchers have identified downflowing counterparts to spicule-related events, features that move toward strong magnetic field regions instead of away from them. A statistical analysis of nearly 20,000 upward-directed events and 15,000 of the redshifted (and newly identified downflowing) counterparts found that these downflows are remarkably similar to their upward-moving siblings in every respect except their direction of motion. This suggests the chromosphere’s jet activity is even more bidirectional than previously appreciated.
Fibrils and the Magnetic Canopy
Spicules aren’t the only thread-like structures in the chromosphere. When you image the Sun in the light of hydrogen-alpha (Hα), a strong red emission line, the chromosphere appears covered in elongated dark and bright streaks called fibrils. These features dominate the appearance of the chromosphere above and around active regions, and their origin has been debated for years.
High-resolution observations from ground-based solar telescopes, combined with numerical simulations, indicate that dynamic fibrils are most likely formed by chromospheric shock waves. Convective flows and global oscillations in the layers below leak upward along the magnetic field lines of concentrated flux regions, steepen into shocks, and drive short-lived jet-like features that trace out the local magnetic geometry. The relationship between fibrils and field lines is not perfectly one-to-one, however. Some fibrils trace a single magnetic field line along their length, while others sample different field lines at different positions. Fibrils tend to follow the same field line for only about 200 seconds, which is shorter than the fibril’s own lifetime. The magnetic landscape underneath is shifting faster than the visible feature can keep up.
The dark appearance of fibrils in Hα images has a specific physical cause. Ridges of enhanced mass density aligned with the magnetic field raise the average height at which the Hα line forms, and because temperature generally drops with height in the chromosphere’s lower layers, these higher-formed regions appear darker against their deeper, brighter surroundings. The width of the Hα line core, meanwhile, is a good thermometer for the gas temperature in the region where the line forms.
Zooming out, the magnetic field itself changes character between the photosphere and the chromosphere. In sunspot regions, the difference in field strength between the photosphere and the upper chromosphere shrinks as you move from the spot’s center toward its edge. Outside the sunspot boundary, the horizontal component of the field is actually stronger in the chromosphere than in the photosphere, a hallmark of what’s called a magnetic canopy: field lines that spread out and become more horizontal as they rise, blanketing the chromosphere above quiet-Sun regions adjacent to strong-field concentrations.
Observing the Chromosphere Through Spectral Lines
The chromosphere is largely invisible in ordinary white light, which is why solar physicists rely on a handful of spectral lines that form preferentially in chromospheric conditions. The most famous is Hα at 656.3 nanometers, which gives the chromosphere its red color during eclipses. Others include the calcium II H and K lines near 393 and 397 nanometers, the magnesium II h and k lines near 280 nanometers, and the calcium infrared triplet line at 854.2 nanometers.
Each line acts as a window into a slightly different slice of the chromosphere, and interpreting what you see through these windows is far from straightforward. Modeling Hα, for instance, requires three-dimensional radiative transfer calculations because the probability of a photon being destroyed (absorbed and converted to heat) is very low in this line. Photons scatter many times before being lost, so conditions far from the point of observation can influence what you measure. Simulations of an active region show that Hα, the Mg II h and k lines, and the Ca II H and K lines can form at surprisingly similar heights, contrary to older assumptions that each line neatly probes a distinct altitude.
A newer addition to the chromospheric toolkit is the Atacama Large Millimeter/submillimeter Array (ALMA), a radio observatory that measures brightness temperatures at millimeter wavelengths. Joint analyses of ALMA data and Hα observations reveal a strong correlation between Hα line width and ALMA brightness temperature. Both diagnostics turn out to share a common physical driver: the population of hydrogen atoms in the first excited energy level. When that population is high, both the millimeter-wave opacity and the Hα opacity increase together, and both diagnostics shift to forming at similar atmospheric heights. The link is mediated by the ultraviolet photon field, which regulates hydrogen’s excitation state throughout the chromosphere.
For tracking chromospheric activity over long timescales, satellite instruments use the Mg II core-to-wing ratio, an index that compares the brightness at the center of the magnesium line to its wings. This index is a reliable proxy for solar ultraviolet variability and correlates strongly with sunspot area. An interesting subtlety is that during solar minimum, when sunspot area drops to zero, the Mg II index does not bottom out completely, because small-scale network magnetic fields scattered across the quiet Sun continue to contribute emission in the Mg II line even when large active regions are absent.
Flares and Chromospheric Evaporation
During a solar flare, enormous amounts of energy are released in the corona when stressed magnetic field lines reconnect. Much of that energy is channeled downward along field lines in the form of accelerated electrons and thermal conduction fronts. When this energy slams into the dense chromosphere, it heats the gas so violently that chromospheric material expands upward into the corona at high speed, a process called chromospheric evaporation. The term is a bit misleading, since no liquid is involved: it refers to dense, cool gas being rapidly heated and launched upward, filling coronal loops with hot, bright plasma that we observe in X-rays.
Observations from the STIX instrument on Solar Orbiter have tracked this process by measuring where hard X-ray sources form at different energies. In the chromosphere, the data show fast upward mass flows reaching speeds of up to 250 kilometers per second, consistent with explosive evaporation. Higher up, in the low corona, a steadier flow of plasma is observed at up to about 100 kilometers per second. The interplay between energy deposition in the chromosphere and mass transport back up into the corona is central to how flares evolve and how their energy is distributed through the solar atmosphere.
Small-Scale Reconnection Events
Flares get the headlines, but much smaller reconnection events happen constantly in and around the chromosphere. Two of the most studied are Ellerman bombs and ultraviolet (UV) bursts. Ellerman bombs appear as small, transient brightenings in the wings of the Hα line and are associated with reconnection at photospheric or low-chromospheric heights, typically where opposite-polarity magnetic fields are pushed together by convective motions. UV bursts, seen in hotter ultraviolet spectral lines, represent reconnection at somewhat greater heights.
When Ellerman bombs and UV bursts appear at the same location and time, they likely represent different layers of a single reconnection system. The reconnection occurs along a vertical or nearly vertical current sheet, with heating and bidirectional jets marking the site. As the reconnection proceeds, its effects are felt at progressively different heights: Ellerman-bomb signatures dominate lower down, while UV-burst signatures appear higher up. This sequential, multi-layered picture of reconnection underscores that the chromosphere isn’t a passive slab of gas; it’s a site of constant, small-scale energy release that may contribute meaningfully to the overall heating budget.
Waves That Thread the Chromosphere Into the Corona
Beyond acoustic waves, magnetoacoustic waves play a particularly important role in sunspot regions. Sunspot umbrae (the dark central areas of sunspots) are known to oscillate with a characteristic three-minute period. These oscillations are driven at photospheric levels and propagate upward along the magnetic field into the chromosphere and beyond. Observations of coronal fan loops rooted in sunspot umbrae show that the three-minute slow magnetoacoustic waves seen in the corona can be traced back to oscillations at the photospheric footpoints of those loops. The chromosphere serves as a waveguide, funneling oscillatory energy along the expanding magnetic field from the dense photosphere into the tenuous corona.
Measuring Chromospheric Magnetic Fields
Measuring magnetic fields in the photosphere is routine: the Zeeman effect splits spectral lines in proportion to the field strength, and the splitting is large enough to detect. In the chromosphere, things are harder. The gas is hotter and more turbulent, which broadens spectral lines and washes out the subtle Zeeman splitting from weaker fields. For fields below a few hundred gauss, which covers most of the quiet chromosphere, the Zeeman effect alone is not enough.
This is where the Hanle effect comes in. The Hanle effect modifies the polarization of spectral-line radiation that has been produced by the scattering of anisotropic sunlight. It is sensitive to much weaker fields than the Zeeman effect, making it essential for probing the quiet chromosphere’s magnetism. The only way to get reliable empirical information on the intensity and topology of these weak fields is through careful measurement and rigorous interpretation of polarization signals in chromospheric lines. Researchers use lines like the calcium I line at 422.7 nanometers (the λ4227 line) to synthesize both Hanle and Zeeman polarization signals, combining forward-scattering geometry with time-dependent three-dimensional magnetohydrodynamic models of the solar atmosphere. The result is a much richer picture of how chromospheric magnetic fields evolve in space and time than the Zeeman effect alone could provide.
Progress in this area has been slow relative to photospheric magnetometry, partly because the required polarimetric sensitivity is extreme. Upcoming solar telescopes with larger apertures and more sensitive polarimeters are expected to push the boundary significantly, but for now, chromospheric magnetic field measurements remain one of the hardest observational challenges in solar physics.
Chromospheres on Other Stars
The Sun is not unique in possessing a chromosphere. Cool stars with outer convective envelopes, roughly those with surface temperatures below about 7,000 K, generally show chromospheric emission in the same spectral lines used to study the Sun’s chromosphere, particularly Ca II H and K and Mg II h and k. The Mount Wilson Observatory’s decades-long monitoring program measured Ca II H and K emission in hundreds of nearby Sun-like stars, revealing that many of them undergo activity cycles analogous to the Sun’s 11-year cycle. The theoretical framework connects these cycles to magnetic dynamos driven by the interplay of differential rotation and convection. Stars that rotate faster tend to have stronger chromospheric emission, because faster rotation amplifies the dynamo that generates magnetic fields, and magnetic fields are the ultimate energy source for chromospheric heating.
Studying stellar chromospheres puts the Sun in context. The Sun turns out to be a relatively quiet star compared to many of its peers, sitting on the lower end of the chromospheric-activity spectrum for stars of its age and spectral type. Younger, faster-spinning stars can have chromospheric emission many times stronger, with correspondingly more frequent and energetic flares. This perspective matters not just for stellar physics but for understanding the space-weather environments of exoplanets orbiting other stars, where chromospheric activity translates directly into ultraviolet radiation levels and particle bombardment.
Why the Chromosphere Is So Difficult to Model
Several features of the chromosphere conspire to make it one of the most computationally demanding regions of the Sun to simulate realistically. The gas is neither fully ionized nor fully neutral, which means that interactions between ions and neutral atoms (a process called ambipolar diffusion, or ion-neutral friction) matter for energy transport and wave damping. The radiation field is far from thermal equilibrium: photons in strong spectral lines scatter dozens of times before being absorbed, so local conditions cannot be determined without knowing the radiation field over a large volume. Three-dimensional radiative transfer is essential for modeling hydrogen lines accurately because of this low photon-destruction probability.1The Astrophysical Journal. THE FORMATION OF THE Hα LINE IN THE SOLAR CHROMOSPHERE
The magnetic field geometry adds another layer of complexity. In the photosphere, the field is largely confined to narrow concentrations between convective cells. By the chromosphere, those field lines have fanned out to fill much of the available volume, creating the magnetic canopy. The transition from a mostly field-free gas to a magnetically dominated one happens within the chromosphere itself, and the height at which it happens varies from place to place depending on the local field strength below. Simulations have to capture this transition accurately or they get the wave propagation, energy deposition, and spectral-line formation wrong.
Recent three-dimensional non-equilibrium magnetohydrodynamic simulations have made impressive strides. One simulation of an active region reproduced long fibrils stretching between opposite-polarity sunspots, reaching up to about 4 megameters in height, and showed that the major chromospheric spectral lines (Hα, Mg II h and k, Ca II H and K) form at similar heights rather than at neatly separated layers.2EDP Sciences (Astronomy & Astrophysics). Three-dimensional modeling of chromospheric spectral lines in a simulated active region That finding has practical consequences: it means observers can’t simply stack these spectral lines into a clean altitude sequence when interpreting images. The chromosphere’s structure is three-dimensional, and the formation heights of its diagnostic lines depend strongly on local conditions.