The Sun is built like a series of nested shells, each with distinct physical conditions and a different way of moving energy outward. At its center sits an extraordinarily dense, hot core where hydrogen fuses into helium, generating nearly all of the Sun’s power. That energy then passes through a radiative zone where photons bounce around for tens of thousands of years, a thin shear layer called the tachocline, a turbulent convective zone where hot gas physically rises and cool gas sinks, and finally the visible surface layers and atmosphere that extend into interplanetary space. Each of these regions behaves so differently from the others that the Sun is less a uniform ball of gas than a stack of radically different environments sharing the same address.
The Core
The innermost roughly quarter of the Sun’s radius is where the action starts. Temperatures exceed 15 million kelvins, and the density is about 150 times that of water. Under these conditions, hydrogen nuclei overcome their mutual electrical repulsion and fuse. The dominant sequence, called the proton-proton chain, begins when two protons collide and merge, releasing a low-energy neutrino in the process. This particular reaction is the first step in a chain that ultimately converts four hydrogen nuclei into one helium nucleus, releasing energy along the way. Spectral detection of the neutrinos from that initial proton-proton step has confirmed that about 99 percent of the Sun’s luminosity comes from this chain.1Nature. Neutrinos from the primary proton–proton fusion process in the Sun2Nature. Comprehensive measurement of pp-chain solar neutrinos
A small fraction of solar energy comes from a secondary cycle involving heavier elements like carbon and nitrogen as catalysts. This pathway matters more in stars heavier than the Sun, but it still leaves traces in our star. The SNO+ experiment recently reported the first evidence of solar neutrinos interacting with carbon-13 nuclei, a detection consistent with theoretical predictions and achieved at a statistical significance of 4.2 sigma.3Phys Rev Lett. First Evidence of Solar Neutrino Interactions on ^{13}C These neutrino measurements do more than confirm textbook physics: they give us a direct probe of conditions deep inside the core, a place no telescope can image.
The Radiative Zone
Energy leaves the core as high-energy photons, but those photons do not fly straight out. Instead, they enter the radiative zone, which stretches from about 0.25 to 0.7 solar radii. Here the gas is still far too hot and dense for convection to take hold, so energy moves by radiation: each photon travels a tiny distance before being absorbed by an atom and re-emitted in a random direction. This random-walk process is astonishingly slow. Standard solar models estimate that the total time for energy to diffuse outward through the Sun is on the order of 170,000 to 200,000 years, a figure that has actually decreased slightly over the Sun’s lifetime as its internal structure has evolved.4Mapana Journal of Science. On the Tachocline Zone Location in the Sun, the Luminosity Transport time scale, the Rotational Inertia and their Time Variation in Standard Solar Evolution Models So the sunlight warming your face left the core during the last ice age, give or take.
Inside the radiative zone, the gas rotates roughly as a solid body, meaning every latitude turns at about the same rate. This uniform rotation is important because it sets the stage for a dramatic change at the zone’s outer boundary.
The Tachocline
Sandwiched between the radiative zone below and the convective zone above is a thin transition layer called the tachocline. It sits near 0.7 solar radii and is thought to be only a few percent of the Sun’s radius thick. What makes it significant is the shear: the solid-body rotation of the radiative interior meets the differential rotation of the convective envelope, where the equator spins faster than the poles. That meeting produces strong gradients in both the radial and latitudinal directions.5PubMed Central. Dynamics of the Tachocline
This shear is widely believed to play a central role in generating the Sun’s large-scale magnetic field. The stretching and winding of magnetic field lines in the tachocline helps build up the toroidal (east-west) magnetic flux that eventually rises through the convective zone and breaks through the surface as sunspot pairs. Observations of bipolar sunspot groups confirm that they originate from toroidal magnetic flux in the convective envelope, and the net toroidal flux generated by differential rotation within each hemisphere can be traced back to the magnetic field distribution at the surface.6PubMed. The crucial role of surface magnetic fields for the solar dynamo In other words, the tachocline is the engine room of the solar dynamo, even though it occupies a sliver of the Sun’s volume.
The Convective Zone
From about 0.7 solar radii outward to the visible surface, energy transport shifts from radiation to convection. The gas here is cool enough (relatively speaking) that it becomes opaque enough to trap photons efficiently, so the temperature gradient steepens and hot parcels of gas physically rise while cooler parcels sink. This churning creates a hierarchy of convective cells visible at the surface.
The smallest and most obvious are granules, bright cells roughly 1,000 kilometers across that tile the entire visible disk and overturn every few minutes. At a larger scale, supergranules span about 30,000 kilometers and live for a day or so. Three-dimensional simulations suggest that supergranulation represents the largest scale at which buoyancy can efficiently drive convection in the Sun’s upper layers. Below that scale, granular downflows dilute as they descend, and the transition from strong buoyant driving to near-adiabatic stratification limits convective power at scales larger than supergranulation.7The Astrophysical Journal Letters. SUPERGRANULATION AS THE LARGEST BUOYANTLY DRIVEN CONVECTIVE SCALE OF THE SUN Whether even larger “giant cells” exist deeper down remains an open question. Some models predict them, but observational evidence has been elusive.
The convective zone is also where the Sun’s differential rotation lives. Gas near the equator completes a rotation in about 25 days, while gas near the poles takes closer to 35 days. This difference, maintained across the full depth of the convective zone, is part of what generates and sustains magnetic activity.
The Photosphere
The photosphere is the thin layer, only a few hundred kilometers deep, from which nearly all the sunlight we see escapes. Its average temperature is around 5,500 kelvins, and it is where the Sun transitions from opaque to transparent. This is the “surface” in any practical sense, even though the Sun has no solid boundary.
The most striking features on the photosphere are sunspots: dark patches where intense magnetic fields suppress convection and lower the local temperature by roughly 1,500 kelvins compared with the surrounding gas. Inside the darkest part of a sunspot, the umbra, convection does not shut down entirely. Realistic three-dimensional simulations show that narrow, nearly field-free upflow plumes punch through the strong magnetic field, creating small bright features called umbral dots. These plumes assume a cusp-like shape near the surface where the rising plasma loses buoyancy, and the resulting dots match observed umbral dots in brightness, size, and lifetime.8The Astrophysical Journal. Magnetoconvection in a Sunspot Umbra So even in the most magnetically dominated parts of the Sun’s surface, convection finds a way.
Beyond sunspots, the photosphere hosts faculae (bright magnetic patches most visible near the limb), pores (small dark spots without a penumbra), and the ever-present granulation pattern driven by convection from below.
The Chromosphere and Spicules
Just above the photosphere lies the chromosphere, a layer roughly 2,000 kilometers thick that is normally invisible against the photosphere’s glare but can be seen during eclipses as a thin reddish ring. Temperatures in the chromosphere actually rise with altitude, from about 4,000 kelvins near the bottom to around 20,000 kelvins near the top, reversing the steady decline seen deeper in the Sun. This marks the beginning of one of the most persistent puzzles in solar physics: how the outer atmosphere gets hotter as you move away from the energy source.
The chromosphere is dominated by spicules, jet-like features that shoot plasma upward at speeds of 20 to 100 kilometers per second and reach heights of several thousand kilometers before falling back or fading. They blanket the Sun’s limb in a constantly shifting forest. New simulations have revealed that many spicules may not be the simple one-dimensional spikes they appear to be. Instead, they can be manifestations of two-dimensional plasma sheets, with the appearance of spinning created by hot swirling plasma columns extending upward to coronal heights.9arXiv. Spinning solar jets explained through the interplay between plasma sheets and vortex columns
What drives spicules has been debated for decades. Numerical simulations have explored whether dumping energy into the mid-chromosphere can accelerate material to the observed heights. The results suggest that pressure gradients from localized heating can propel plasma upward, but the process is far less efficient than once thought: most of the deposited energy gets radiated away, and hydrogen ionization absorbs much of the rest, limiting the thermal energy available to push gas higher.10The Astrophysical Journal. NUMERICAL SIMULATIONS OF SPICULE ACCELERATION Magnetic mechanisms, including tension in twisted field lines, are now considered more promising drivers for the tallest and fastest spicules.
The Corona and the Heating Problem
Above the chromosphere, separated by a narrow transition region where temperatures skyrocket over just a few hundred kilometers, lies the corona. This tenuous outermost atmosphere stretches millions of kilometers into space and reaches temperatures of one to three million kelvins. The corona is the source of the solar wind and the stage for some of the Sun’s most violent events, including flares and coronal mass ejections.
The central mystery of the corona is its temperature. The photosphere sits at roughly 5,500 kelvins; the corona is hundreds of times hotter. Heat does not normally flow from a cooler source to a warmer destination without some extra mechanism at work. Two broad families of explanation have dominated the debate: wave heating and magnetic reconnection.
In the wave-heating picture, magnetohydrodynamic waves, particularly Alfvén waves, are generated lower in the atmosphere and carry energy into the corona, where they dissipate and heat the gas. Magnetic reconnection in spicules has been proposed as one source of these waves.11Iranian Journal of Astronomy and Astrophysics. Generation of Alfvén Waves by Small-Scale Magnetic Reconnection in Solar Spicules In the reconnection-heating picture, tangled and twisted magnetic field lines in coronal loops snap and reconnect, converting magnetic energy directly into heat. Simulations of kink-unstable loops show that reconnection in the nonlinear phase can efficiently dissipate magnetic energy, heating the loop to high temperatures throughout its volume.12Astronomy & Astrophysics. Coronal heating by magnetic reconnection in loops with zero net current The current consensus is that both mechanisms contribute, with their relative importance varying between different types of coronal structures.
The Solar Wind and Heliosphere
The corona does not simply sit above the Sun; it continuously expands outward as the solar wind, a stream of charged particles that fills the entire solar system. The wind comes in distinct flavors. Fast solar wind, blowing at 600 to 800 kilometers per second, originates mainly from coronal holes, regions of open magnetic field near the poles. Slow solar wind, at roughly 300 to 500 kilometers per second, tends to come from near the edges of coronal streamers and the boundaries of active regions.
A newer category, Alfvénic slow wind, blurs the traditional fast-slow distinction. Despite traveling at slow-wind speeds, it carries large-amplitude Alfvén waves and has composition resembling fast wind. Analysis integrating remote and in-situ measurements indicates that Alfvénic slow wind behaves like fast wind in terms of its acceleration: a wave-pressure gradient is needed to explain its full speed, whereas non-Alfvénic slow wind can be driven simply by thermal pressure gradients from its electrons and protons.13The Astrophysical Journal. Differentiating the Acceleration Mechanisms in the Slow and Alfvénic Slow Solar Wind The implication is that the traditional two-type classification of solar wind is too simple; the acceleration physics depends on the magnetic topology at the source.
The solar wind inflates a vast bubble around the Sun called the heliosphere, which extends well past the orbit of Pluto. At its outer boundary, the wind’s pressure balances the pressure of the interstellar medium. Voyager 1 and 2 crossed this boundary and confirmed its existence, but the shape and dynamics of the heliosphere remain active research topics.
Magnetic Switchbacks and What Parker Solar Probe Found
NASA’s Parker Solar Probe, launched in 2018, has been flying closer to the Sun than any previous spacecraft, dipping into the corona itself. Among its most surprising discoveries are magnetic switchbacks: sudden, sharp reversals in the direction of the solar wind’s magnetic field. These S-shaped kinks in the field had been seen before at greater distances, but Parker found them to be far more common and dramatic close to the Sun.
An analysis of the probe’s first eight encounters identified 501 robust switchbacks. More than 46 percent of them were classified as rotational discontinuities, in which the magnetic field rotates smoothly across the boundary, while about 6 percent were tangential discontinuities.14The Astrophysical Journal Letters. Magnetic Switchbacks Heat the Solar Corona A broader study cataloging more than 1,000 switchbacks across seven encounters found many more current sheets inside switchbacks than outside them, suggesting that these internal structures help stabilize the S-shaped configuration.15The Astrophysical Journal. The Structure and Origin of Switchbacks: Parker Solar Probe Observations
The origin of switchbacks is still debated. Some models tie them to interchange reconnection at the Sun’s surface, where open and closed field lines swap footpoints. Others suggest they form in the solar wind itself through velocity shear or turbulence. Whatever their origin, switchbacks carry energy, and their dissipation may contribute to heating and accelerating the solar wind, linking them back to the coronal heating problem.
The Solar Activity Cycle and Coronal Mass Ejections
The Sun’s magnetic field does not stay constant. It follows an approximately 11-year cycle during which sunspot numbers rise and fall, the large-scale field flips polarity, and the frequency of flares and coronal mass ejections (CMEs) waxes and wanes. This cycle is driven by the solar dynamo, a feedback loop in which differential rotation and convective motions stretch, twist, and amplify magnetic fields. As noted earlier, the tachocline plays a key role by generating toroidal flux from the roughly dipolar surface field that peaks around activity minima.16PubMed. The crucial role of surface magnetic fields for the solar dynamo
Coronal mass ejections are enormous eruptions of magnetized plasma from the corona, carrying billions of tons of material into interplanetary space at speeds ranging from a few hundred to over 2,000 kilometers per second. The fastest CMEs appear to be driven primarily by magnetic reconnection in the flare beneath them, while slower events may depend more on other mechanisms such as gradual magnetic pressure buildup.17The Astrophysical Journal. How Does Magnetic Reconnection Drive the Early-stage Evolution of Coronal Mass Ejections? When aimed at Earth, CMEs can trigger geomagnetic storms that disrupt power grids, satellite communications, and navigation systems, making the Sun’s structure and activity a matter of practical concern.
The Solar Abundance Problem
Building a computer model of the Sun’s interior requires knowing its chemical composition, and this is where things get contentious. Solar models are calibrated against two independent sets of observations: helioseismology (using sound waves that propagate through the interior to map its structure) and neutrino fluxes from the core. Both of these constrain properties like the depth of the convective zone and the helium abundance at the surface.18The Astrophysical Journal. THE CHEMICAL COMPOSITION OF THE SUN FROM HELIOSEISMIC AND SOLAR NEUTRINO DATA
In the early 2000s, revised estimates of the Sun’s heavy-element abundances (metals, in astronomer-speak) dropped significantly below older values. Models built with these lower metallicities stopped agreeing with helioseismic data, particularly the measured depth of the convective zone and the sound-speed profile. This became known as the “solar abundance problem.” A natural hope was that raising the metal abundances back up might fix things. Recent work, however, shows that when high-metallicity models are forced to also reproduce the Sun’s observed lithium depletion, new tensions appear with helioseismology and neutrino data.19Astronomy & Astrophysics. Higher metal abundances do not solve the solar problem The problem, in other words, is not just about what the Sun is made of. Something in the physics of the models themselves, whether it involves opacity, mixing, or diffusion, needs to be improved.
Tracking Solar Activity Over Millennia
Direct sunspot records go back only about four centuries, to the early 1600s. To understand solar variability on longer timescales, scientists turn to cosmogenic isotopes: radioactive atoms like beryllium-10 and carbon-14 that are produced when cosmic rays strike atoms in Earth’s atmosphere. When the Sun is more active, its stronger magnetic field deflects more cosmic rays, so fewer cosmogenic isotopes are produced. By measuring these isotopes in ice cores and tree rings, researchers can reconstruct solar activity stretching back roughly 12,000 years, covering the entire Holocene epoch.20PubMed Central. 9,400 years of cosmic radiation and solar activity from ice cores and tree rings21Living Reviews in Solar Physics. A history of solar activity over millennia
These records reveal that the Sun has gone through extended periods of unusually low activity, called grand minima, as well as bursts of elevated activity. The Maunder Minimum, spanning roughly 1645 to 1715, is the best-known example of a quiet spell, during which sunspots nearly vanished. The proxy record shows that such episodes have recurred multiple times over the past several thousand years. Understanding whether the Sun’s internal structure subtly shifts during these quiet phases, perhaps through changes in the dynamo or the tachocline’s behavior, is an active area of research.
How the Sun’s Structure Compares to Other Stars
The Sun’s layered interior, with a radiative core wrapped in a convective envelope, is typical for stars of similar mass. But this architecture is not universal. Stars significantly more massive than the Sun have convective cores and radiative envelopes, essentially the reverse arrangement, because the nuclear reactions in their cores are more temperature-sensitive and generate steeper energy gradients. Meanwhile, the smallest and coolest stars, the M dwarfs, can be fully convective from core to surface, with no radiative zone at all. Somewhere in the M-dwarf mass range, the radiative interior disappears entirely and the whole star churns.
This matters because the tachocline, the shear layer between the radiative and convective zones, is thought to be critical for generating a Sun-like magnetic cycle. Fully convective stars lack a tachocline, yet many of them are magnetically active, sometimes violently so. How they generate and sustain their magnetic fields without that shear interface is one of the open questions in stellar physics. It highlights that the Sun’s particular layered structure is not the only way a star can build a magnetic field; it is just the way that happens to work at our star’s mass and age.