Sol — Our nearest star

Sol commentary

Sol

The Sun’s gravitational attraction holds the planets in their orbits, and its heat and light make life possible on Earth. Its diameter (864,000 miles) implies a volume capable of holding more than a million Earths, and its mass accounts for about 99.85% of the mass in the Solar System.

It’s a fairly typical star, but uniquely important to us because it’s close: about 150 million km away. Solar energy provides essentially all the heat and light received by Earth, sustaining every living organism. Solar radiation and the solar wind also shape our near-space environment, influencing Earth’s magnetosphere and ionosphere. Historically, studying the Sun has driven advances in atomic and nuclear physics, plasma physics, and magnetohydrodynamics — and modern satellites have enabled direct investigation of how solar activity couples into Earth’s magnetic field.

Solar structure

The Sun formed from a collapsing cloud of gas and dust about 4.6 billion years ago. Its mass (≈330,000 Earth masses) generates central pressures and temperatures high enough for nuclear fusion. A small central region — the core — contains a large fraction of the mass and is responsible for essentially the entire luminosity.

Inside the Sun

Solar power is produced mainly by the proton–proton chain, converting hydrogen into helium. Fusion releases energy by a tiny mass deficit (mass → energy). Over billions of years, the core’s composition evolves: hydrogen falls and helium rises, while the Sun’s total mass changes only slightly during its main-sequence lifetime. On long time scales, depletion of central hydrogen will force structural changes, leading to a red giant phase and ultimately a white dwarf.

The physical conditions driving these reactions cannot be measured directly. We cannot put a thermometer in the Sun's core; instead, solar models are constrained by observations of luminosity, spectrum, neutrinos and solar oscillations. Those models put the central temperature at about 1.5 × 10⁷ K, falling steadily with increasing radius to about 5800 K at the visible surface.

Energy generated in the core does not stream directly outward. The solar interior is effectively opaque to electromagnetic radiation: photons are repeatedly absorbed, re-emitted and scattered, so energy works its way through the radiative zone by a very slow random walk. Estimates depend on exactly what is being counted, but a commonly quoted figure is roughly 170,000 years for radiative energy to reach the top of the convection zone. From there convection carries energy onward toward the photosphere. The important point is that energy made in the core does not simply fly straight out in eight minutes.

The density structure of the Sun is equally extreme. Although the average density of the Sun is only about 1.4 times that of water, the density rises sharply toward the centre. At the core it reaches roughly 1.6 × 10⁵ kg m⁻³, comparable to twelve times the density of lead. This strong central concentration means that a large fraction of the Sun’s mass is confined to its inner regions: approximately 90% of the total mass lies within about half the solar radius.

These extreme temperatures and densities ensure that matter in the core is fully ionised. Atoms cannot retain electrons, and the plasma consists primarily of protons, helium nuclei, and free electrons. Under these conditions, nuclear fusion proceeds steadily, powering the Sun and maintaining its remarkably stable output over billions of years.

Outside the core, the radiative zone extends to roughly 0.7 solar radii. Above that, through the outermost ~200,000 km of the interior, convection becomes the more efficient way to move energy: hot plasma rises, cools and sinks again. The convection zone reaches all the way to the visible surface, where its smallest readily visible cells appear as granulation.

Convection is often described as layered in scale: very large “giant cells”, intermediate “supergranules” tens of thousands of km across, and smaller granules roughly ~1,000 km across (and ~minutes in lifetime) that form the photospheric texture.

Fusion also produces neutrinos that interact so weakly with matter that they escape directly. Early neutrino experiments famously observed fewer solar neutrinos than predicted, a puzzle that helped drive refinements in both solar modelling and particle physics.

The visible surface

The bright “surface” is the photosphere — the sharp solar disc seen visually (with safe methods). Under good conditions it appears mottled: granulation. Bright granule centres correspond to rising hot plasma, while darker lanes mark cooler sinking flows. The pattern changes continuously as granules dissolve and reform, so the photosphere can look noticeably different over minutes.

The limb appears surprisingly sharp rather than gradually fading. That tells us most visible light comes from a relatively thin layer compared with the Sun’s radius — the atmosphere transitions quickly from opaque to transparent at visible wavelengths as density and ion chemistry drop.

The disc is also dimmer toward the edge (limb darkening). Looking near the centre we see deeper, hotter layers; near the limb we see higher, cooler layers along a longer, more oblique path. The effect depends on wavelength and is typically stronger in the blue.

You’ll see multiple “surface temperature” numbers quoted because they’re defined in different ways. A blackbody fit to the continuum suggests a value around 6000 K; an effective temperature tied to luminosity is commonly quoted around ~5800 K. The difference reflects that the light we receive is produced across a thin layer with a temperature gradient, and different lines of sight sample different depths.

Beyond the photosphere

The photosphere is not the end of the Sun. Above it lie increasingly thin layers of plasma which are difficult to see in ordinary visible light but dominate images made in hydrogen-alpha, ultraviolet, extreme ultraviolet and X-rays. This is where the Sun stops looking like a fairly tidy glowing ball and starts behaving much more obviously like a magnetic star.

Immediately above the photosphere is the chromosphere. Its name comes from the reddish colour seen briefly around the edge of the Moon during a total solar eclipse. Density continues to fall with height, but the temperature, after reaching a minimum above the photosphere, begins to rise again. The chromosphere is threaded by magnetic fields and is full of short-lived jets, spicules and other moving structures. In hydrogen-alpha light it can look dramatically different from the calm white-light disc.

Above that is the transition region, a thin and irregular boundary where temperatures climb extremely rapidly from tens of thousands to hundreds of thousands of kelvin. There is no neat spherical shell that could sensibly be painted onto the 3D model above; magnetic loops, flows and local structure make the real boundary decidedly untidy.

The outer atmosphere is the corona. It is extraordinarily tenuous, yet typically reaches temperatures of around one to two million kelvin and can be even hotter in active regions. That is hundreds of times hotter than the photosphere below it. Explaining exactly how the corona is heated remains an active problem in solar physics, with magnetic reconnection and waves among the important parts of the story.

During a total solar eclipse the corona becomes visible to the eye as a pale halo extending far beyond the darkened Sun. Space telescopes and coronagraphs can study it without waiting for the Moon to get in the way, and ultraviolet and X-ray observations reveal loops and active regions that are almost invisible in ordinary white light.

A magnetic and rotating Sun

The Sun is plasma rather than a solid body, so it does not rotate as one rigid object. The equatorial regions turn in about 25 days while higher latitudes rotate more slowly. Together with convection inside the Sun, this differential rotation helps twist and reorganise the solar magnetic field. Much of what we describe as solar activity is ultimately magnetic activity.

Sunspots are the most obvious visible sign. They are regions where strong magnetic fields suppress some of the normal convective flow, leaving the photosphere locally cooler. They only look black by comparison with the much hotter material around them; taken on their own they would still be very bright indeed. Large spot groups can be bigger than Earth.

The number of sunspots rises and falls in a cycle averaging about 11 years. Near solar minimum the disc can remain almost spotless for long periods; around solar maximum active regions are much more common. The magnetic polarity of the Sun reverses during each 11-year cycle, so returning to approximately the same global magnetic orientation takes about 22 years. The cycle is real, but it is not a metronome: individual cycles vary in length and strength.

Flares, prominences and coronal mass ejections

Magnetic fields above active regions can store enormous amounts of energy. When their configuration changes rapidly, particularly through magnetic reconnection, that energy can be released in a solar flare. A flare is primarily a burst of electromagnetic radiation and energetic particles. Its X-rays and extreme ultraviolet radiation travel at the speed of light, so if the eruption is on the Earth-facing side we see its radiative effects about eight minutes after they leave the Sun.

A prominence is comparatively cool, dense plasma suspended above the photosphere by magnetic fields. Seen against space at the limb it can form a huge glowing arch; seen projected against the bright disc in hydrogen-alpha it appears darker and is usually called a filament. Some remain stable for days or weeks, while others erupt.

A coronal mass ejection, or CME, is something different again: a vast cloud of magnetised plasma thrown out through the corona into interplanetary space. Flares and CMEs often occur together but one does not guarantee the other. An Earth-directed CME commonly takes of order one to three days to arrive, depending on its speed and on conditions in the solar wind ahead of it.

The solar wind and the heliosphere

The corona does not have a sharp outer edge. Some of its plasma is continuously accelerated away from the Sun as the solar wind, carrying the Sun's magnetic field with it. Near Earth, typical solar-wind speeds are a few hundred kilometres per second, with faster streams commonly associated with coronal holes. The wind is thin enough to count as an excellent laboratory vacuum, but over astronomical distances it matters enormously.

The solar wind inflates the heliosphere, the enormous bubble of solar plasma and magnetic influence surrounding the planets. Its outer boundary, the heliopause, lies far beyond Neptune. In that sense Earth, the other planets and most familiar members of the Solar System live inside an extended atmosphere of the Sun. Voyager 1 and Voyager 2 are the only spacecraft so far to have crossed the heliopause and sampled interstellar space directly.

From the Sun to space weather

Solar activity matters to Earth because our magnetic field and upper atmosphere respond to changes in the solar wind, energetic particles and solar radiation. The effects are collectively called space weather. Most of the time the consequences are modest, but stronger events can produce aurora far from the polar regions and disturb HF radio propagation, satellite operations, navigation systems and, in extreme cases, electrical infrastructure.

The Kp value shown near the top of this page belongs on the Earth end of that chain. It describes how disturbed Earth's magnetic field is on a planetary scale; it does not tell us by itself what the Sun has just done. A high Kp can be the result of a CME, a high-speed solar-wind stream or other changes in the interplanetary magnetic field. For radio purposes this distinction matters: the same solar event can improve some propagation paths, wreck others, and produce a very pretty aurora while doing so.

The Sun's future

The Sun is about halfway through its main-sequence life. It has been fusing hydrogen in its core for roughly 4.6 billion years and should continue doing so for about another five billion years. Its output is not perfectly constant over geological time: as helium accumulates in the core the Sun gradually becomes brighter.

Eventually the hydrogen available for fusion in the core will be depleted. The core will contract and heat while the outer layers expand, turning the Sun into a red giant. The exact fate of Earth during that phase is a more complicated question than the usual diagram suggests, but the inner Solar System will certainly become an extremely inhospitable place.

The Sun is not massive enough to end as a supernova. After shedding its outer layers it will leave behind a hot, dense white dwarf, roughly Earth-sized but containing a large fraction of the Sun's present mass. With no further sustained fusion, that remnant will then cool extraordinarily slowly. Fortunately, none of this need interfere with plans for next weekend.