Sol (The Sun)
Above: a photo of the Sun's chromosphere taken with NASA's SOHO observatory. Image courtesy of
SOHO NASA. Click the image to view an animated version.
Star Sol

General Characteristics

Spectral Class: G2V (a yellow dwarf variable, a main sequence star)

Surface temperature (photosphere): 5800 K
Core temperature: 8 000 000 K

Mean distance from the Earth: 1 AU (AU, astronomical unit = 149 598 000 km).
Approximate age: 5 thousand million years.

Radius: 696 000 km (109 times the radius of the Earth).
Mass: 1.99 x 10^30 kg (333 000 times the mass of the Earth).
Luminosity (rate of energy radiation): 3.86 x 10^26 W.
Mean density: 1400 kg/m3 (1.4 times that of water).

The part of Sol visible by ordinary optical means is its atmosphere. The atmosphere is divided into the
innermost visible layer, the yellow photosphere, surrounded by the red chromosphere, which is
surrounded by the tenuous white corona (crown). The Sun, including its atmosphere, is made of very
plasma (a plasma is a highly ionised gas, in which the high temperatures have stripped many of
the electrons off the atoms, forming a gas of electrons intermixed with a gas of ions). Plasma's are
overall neutral, but their individual particles are electrically charged and so moving plasmas generate
magnetic fields and magnetic fields in turn can channel the flow of plasma. Thus, many of the
atmospheric phenomena of the Sun are due to magnetic forces.

The Photosphere

The photosphere (PS) is the layer of Sol's atmosphere that you will normally see should you glance at
the Sun (be very careful when looking at the Sun with your eyes, as the UV light may damage your
retina and NEVER observe the Sun directly through binoculars or a telescope as this will likely destroy
retinal cells! That said, I think everyone has at least glanced casually at the Sun, or even looked at it
for a second or two when it is low on the horizon (and thus dimmest) without causing permanent
damage to their eyes).

The photosphere is that layer of Sol's atmosphere from which photons of visible light escape to be
radiated into space, and so it is the most optically visible layer. The less brilliant outer layers of the
atmosphere, the chromosphere and corona, are only visible to the eye during a total eclipse, but can
be observed with special instruments. However, we can see no deeper into the Sun than the
photosphere, it is the layer at which the Sun's atmosphere becomes opaque to visible light deeper in
and largely transparent to visible light further out.
Left: an image of Sol's photosphere (image courtesy of
SOHO NASA). Click to enlarge. If you enlarge this image then
you can make out faint speckles on the surface, these are
granules, they are bright irregular speckles surrounded by
darker lanes. Also visible in this image are darker
These are cooler regions of the photosphere, and though
still immensely hot, they appear dark in comparison to the
bright surrounds of the hotter regions. Also visible is a
phenomenon called
limb darkening - the edge (limb) of the
Solar disc appear darker. This is because we can see
deeper into the hotter layers of the photosphere when we
look directly at the centre of the disc, but we see only
shallower, cooler and less radiant layers when we look at the
edge. This is because we can only see a fixed distance
through the photosphere and when we look at the limb we
are looking at an angle and so have more gas to see
through to reach the hotter inner layers.
The photosphere is hundreds of km deep. The innermost layers are at a temperature of about 6000K
and the topmost layers at 4000K and so the uppermost layers are less radiant and darker. The
granules have an average diameter of 700-1200 km and an average lifetime ranging from 5 to tens of
minutes. This granulation is the top of Sol's convective zone (CZ). Thew convective zone occurs just
beneath the base of the photosphere and spans 20% of Sol's radius.

Granulation and Convection

Just beneath the photosphere is the convective zone. In this zone heat is largely transported from
the inside-out by convection. Large cells of hotter plasma ascend, cool and then descend as cooler
plasma (hotter gases and plasmas are less dense and so are buoyed upwards). This is an efficient
means of transporting heat. The location and extent of convective layers differ in different star types.
Beneath the convective layer is the
radiative zone, which transports heat from the core to the
convective zone largely by radiation, because at the hotter internal temperatures, the plasma in the
radiative zone is transparent to solar radiation. The core is the hottest region in which radiation is
produced by nuclear fusion takes place, generating enough radiation to prevent the Sun from
collapsing under its own gravity. In the photosphere, the convection manifests itself as visible
convection cells called
granules on the surface.

Sunspots and Solar Cycles

Sunspots are about 3800-3900 K and so appear dark compared to the surrounding photosphere
which is at 5800 K. The central darker region of a sunspot is called the
umbra which is usually
surrounded by a lighter
penumbra. The penumbra is composed of radiating filaments. The largest
sunspots may be more than 60 000 km in penumbral diameter. Sunspots occur in regions where the
magnetic field is about 4 times stronger (0.4 T) than average (0.1 T). These locally strong
magnetic fields probably inhibit convection and so less heat is transported to the surface at sunspots
and hence they are cooler. An alternative/additional explanation is that the strong magnetic fields push
aside the hotter photosphere, exposing the cooler and darker regions of plasma below, giving rise to
the cooler sunspots. Plasma moves outwards away from the umbra in upper filaments and towards the
umbra in lower filaments (at speeds of about 6 km/s).

Sunspots have an associated
magnetic polarity. Magnetic fields always join two magnetic poles
(excluding recent research which has possibly found the long sought after magnetic monopole) - a
north magnetic pole and a south magnetic pole. This terminology has nothing to do with geographic or
geometric poles, except that at present the magnetic north pole of the Earth is close to its geometric
north pole. Magnetic field lines diverge from the magnetic north pole and converge at the magnetic
south pole (magnetic field lines always run from a magnetic north to magnetic south). Stars, like
planets, have magnetic fields resembling those of giant bar magnets. However, small local magnetic
north poles may form elsewhere on the Sun's surface as field lines break out from the surface where
you might otherwise not expect them to (magnetic field lines are a bit like lots of coiled wires and
sometimes they spring out from the sides!).

Usually sunspots are found in magnetically linked pairs - one will be a magnetic north pole from which
magnetic forces and field lines emerge from the Solar surface, whilst the other will be the
corresponding magnetic south pole into which the field lines return into the photosphere. Sometimes,
however, a single sunspot may occur, if one of the magnetic poles is too diffuse to form spots and on
other occasions, one magnetic pole may connect to several complementary magnetic poles in a
sunspot group.

Sunspot numbers cyclically, with a maximum (or minimum) in numbers occurring every 11 years on
average, though the cycle may vary by 2 or 3 years from one cycle to the next. There is also historical
evidence to suggest that the sunspot cycle may pause at a prolonged minimum for several decades.
By convention, a new cycle is said to begin when the sunspot numbers dip at a minimum. Few
sunspots are found near the poles (most occur along a band 40 degrees either side of the Solar
equator) and they tend to appear closer and closer to the equator as the cycle proceeds (appearing
closest to the equator just before a minimum).

The Sun rotates eastward (counterclockwise), like the Earth, and so sunspots appear to move
westward as seen from Earth (but not as seen from space). In magnetically linked pairs, one sunspot is
generally westward of its partner and so apparently moving in front of it and is called the preceding
sunspot whilst its partner is the following sunspot. During one 11-year cycle, all the preceding spots
have the same magnetic polarity (north or south) in the Sun's Northern Hemisphere, but the opposite
polarity in the Sun's Southern Hemisphere. This pattern reverses between the hemispheres in the next
11-year cycle, so that really there is a 22-year cycle of sunspot activity.

Sunspot number increases as solar activity itself increases during the cycle. A large active region,
several hundred thousands of km across, in which magnetic forces are concentrated, is associated
with each sunspot group. This locally high magnetic activity results in several phenomena that are
seen in the overlying chromosphere and corona.
Above: a composite diagram of Solar structure (made with photos courtesy of SOHO/NASA and a Pov-Ray
representation of the interior structure. The labels are as follows: C: core; CME: coronal mass ejection; Co:
corona; CoSt: coronal streamers; CS: chromosphere; CZ: convective zone; F: filament; Fl: flare; G: granules;
P: plages (bright spots);  Pr: prominence; RZ: radiative zone; SG: supergranules; Sp: sprites; SS: sunspots;
SW: solar wind.
Solar Rotation

The Sun exhibits differential rotation - the equator rotates once every 24.47 days and rotates faster
than higher latitude regions. At the start of a solar cycle the magnetic field is weak and dipolar - like a
normal bar magnet one pole of the Sun is magnetic north and the other is magnetic south. The
magnetic field lines run about 70 000 km beneath the surface and mostly emerge at 60 degrees
latitude north and south of the equator. Differential rotation soon complicates this pattern, since the
equator rotates faster, and since plasma can drag magnetic fields with it, the field lines get stretched
at the equator and become tightly wound up. Rising convection cells also move the field lines around
by rotating them locally around their axis until they are wound-up into helices like ropes. Kinks develop
(as they do in a piece of string twisted from both ends) and rise up as bipolar regions - magnetic loops
with one end a magnetic north and the other a magnetic south, generating sunspots. Migration of the
magnetic regions towards the poles neutralises and then reverses the magnetic polarity of the Sun,
establishing the second 11-year phase of the 22-year cycle.

The Core of the Sun

The Sun generates its energy by nuclear fusion processes that take place in its core - the region
where temperatures are hot enough for fusion to take place at a sufficient rate. The core is about 400
000 km in diameter and has a temperature of about 15 million K! At these temperatures, the matter is
completely ionised into a total plasma of atomic nuclei and electrons. The core is also very dense,
containing some 60% of the Sun's mass in only 2% of the Sun's volume. Nuclear fusion in the Sun
proceeds by a chain of reactions called the
p-p chain (proton-proton chain) which generates over
98% of the Sun's energy with 1.6% be generated by a different series of reactions called the
(carbon-nitrogen-oxygen cycle) which becomes much more important in heavier stars than the

For a star to sustain itself by nuclear fusion its core must reach a temperature of at least 10 million K.
It is a common misunderstanding that nuclear fusion makes the Sun hot. Yes, it generates vast
amounts of heat, but only enough to replace what is radiated out into space. The Sun is hot because it
was born from a contracting cloud of cool gas, which warmed up as it contracted, as gravitational
potential energy was converted into thermal energy. Eventually, the core became hot enough to
trigger nuclear fusion and this generated enough pressure to halt further collapse and hence further
heating - nuclear fusion actually stops the Sun collapsing and hence stops it getting hotter!

The (p-p) fusion process begins when two protons, which are essentially hydrogen atomic nuclei
(stripped of their electrons in the hot plasma) collide with sufficient force to make them fuse into a
deuterium nucleus:
The Chromosphere

The layer of the Sun's atmosphere above the photosphere is the chromosphere (lit. 'Coloured
sphere'); so-named because it is red in colour due to hydrogen H-alpha line emission. [The H-alpha
emission line is due to an electron 'jumping' from the n=3 hydrogen atomic orbital down to the n=2
orbit, where n is the principal quantum number and in losing energy the electron emits a photon of red
light at a wavelength of 656.4 nm). The chromosphere is about 2000 km thick.

The chromosphere is less dense than the photosphere and also much less dim and so is not normally
seen as the light from the underlying photosphere shines straight through it. However, it can be briefly
seen as a thin red layer around the Sun's disc during a total eclipse, or through special optics that
filter through only the hydrogen-alpha red light.

The chromosphere has larger convection cells than the photosphere, and in the chromosphere these
are called
supergranules (as opposed to granules in the photosphere) which average about 300
000 km in diameter and have a lifetime averaging about 24 hours. In the supergranules hot plasma
rises at about 50-100 m/s.

Spicules are erupting jets of plasma up to 7000 km high and which each last for 5-10 minutes. In
these jets, plasma shoots up from the photosphere/chromosphere at 20-30 km/s! The spicules usually
occur on the boundaries of supergranules, and so form a network pattern of upward-pointing spikes
across the chromosphere.

The chromosphere has bright
mottles, called plages, which are magnetically active regions associated
with bright patches in the upper photosphere called
faculae. Faculae are hotter regions with
enhanced magnetic fields perpendicular to the Sun's surface which are associated with sunspots -
forming just before sunspot clumps appear and then persisting several weeks after the sunspots
disappear. Chromospheric
fibrils are active regions, about 1100 km long, 725-2200 km wide that each
individually last only 10-20 minutes, but form a pattern persisting for several hours.

prominences are arcs of plasma that follow magnetic field loops in the corona. These
can last for up to about a year before breaking-up and either disappearing or erupting into space. A
quiescent prominence is about 200 000 km long, 40 000 km tall and 600 km thick. In contrast, an
active prominence is smaller (about 60 000 km long) and less stable and appear as long loops, puffs
or sprays of plasma.
Filaments are dark ribbons seen against the disc of the chromosphere and are
simply prominences seen from above.
The corona is characterised by the emission of radio waves. These are due to so-called free-free
transitions, meaning that free electrons in the plasma, not bound to atomic nuclei, lose energy and
emit radio-wave photons in the process. This happens when an electron is scattered by an ion (in a
sense you can think of the electron bouncing off the ion, however, it is perhaps more useful, though
still not accurate, to think of a negatively charged electron passing by a positively charged electron
and being attracted to it causing the electron to alter its path of motion). Higher density regions of the
corona emit shorter wavelength radio-waves such that waves of wavelength 1-20 cm are emitted from
the denser lower corona and chromosphere, whilst wavelengths over 10 cm usually come from the
less dense outer corona.

The corona is extremely hot (counter-intuitive as one might expect the outermost layer of the Sun to
be the coolest) with green spectral emission lines due to Fe XIV (iron ions carrying +13 charge having
been stripped of 13 of their electrons by the immense temperatures) indicating a temperature of 2.3
million K! The Fe XIV line (green light, wavelength = 530.3 nm) and also the red Fe X line (637.4 nm)
forbidden spectral lines. These are emission lines not normally seen, since it takes a
comparatively long time for an electron to make such downward jumps in forbidden spectral lines
(emission lines involve electrons descending energy levels in the atom/ion, emitting the energy they
lose as photons of electromagnetic energy). Normally, atoms/ions in these excited states collide with
other particles and lose their energy as thermal heat instead and so the forbidden emission lines are
not observed (ordinary emission lines that are not forbidden will still be observed as these involve
short-lived excited states in which the electron rapidly jumps down). However, in the thin low density
corona, such collisions are relatively few and far between and this allows sufficient time for the
electrons to jump down.

When a magnetic field line arcs away from the Sun's surface as a
coronal loop, plasma is drawn
along by the magnetic field, forming a chromospheric prominence/filament. When the loop reconnects
to another part of the Sun's magnetic field, field lines from different regions are broken and joined
together with field lines from other regions. (This is thought to occur between field lines in the loop
and those that surround and extend into space). This reconnection generates an immense electric
currents that flow to try and oppose the magnetic change. This releases colossal amounts of energy
(the electrical resistance of the plasma heats it up by dissipating the electric current) resulting in
flares, coronal loops, prominences/filaments and coronal mass ejections. [This can be thought of as
field lines getting bent too far or tangled up until they break and release the 'magnetic strain'].

Solar flares are massive explosions in the Sun's atmosphere that brighten in a few minutes and then
fade over the course of about an hour. They involve all layers of the atmosphere (photosphere,
chromosphere and corona). Flares have prominent Ca XV spectral lines (emission lines due to the
presence of very hot and highly ionised Ca^14+ ions) indicating a temperature of around 3.6 million
K. A typical flare releases about 6 x 10^26 J and ejects plasma at speeds close to the speed of light.
Flares emit a burst of electromagnetic radiation, in the forms of visible light, radio waves and gamma
waves. At the site of reconnection, it is though that a helix of unconnected magnetic field lines
radiates from the Sun, carrying plasma with it in a coronal mass ejection.

Coronal mass ejections (CMEs) eject huge amounts of plasma from the Sun into interplanetary
space, often as an expanding 'bubble'. Up to 1-10 thousand million tonnes of plasma are ejected at
20-3200 km/s. These ejections are associated with the presence of quiescent filaments (they appear
to be prominences that break) and occur from 1-6 times an Earth day.

In more active periods, the corona extends more-or-less uniform around the Solar disc, but at less
active times it is generally confined to the equatorial region, with
coronal holes over the poles.
Coronal streamers are associated with active regions.

Solar tornadoes are plasma jets that eject from the Sun's poles and have 15 km/h wind speeds and
gusts at up to 500 000 km/h and are large enough to swallow the Earth.

Open flux-lines are magnetic field lines that do not arc back down to the Sun's surface, as do coronal
loops, and result in
coronal holes (especially present at the Sun's poles) and the solar wind. The
solar wind is a continuous stream of charged particles emitted by the Sun. They radiate from the Sun
more-or-less along radial lines. In these regions, magnetic field lines accelerate the particles and,
combined with their high thermal energy, causes the particles to reach escape velocity and escape
from the Sun's gravitational field. In the region of the Earth, particles in the solar wind are moving at
about 200-900 km/s. Coronal mass ejection and other eruptive phenomena also contribute to the
solar wind. The Sun's magnetic field is dragged along by the expanding plasma wind (not only can
magnetic fields drag slow-moving plasmas along, but very fast plasmas can drag magnetic fields with
them). This becomes the
interplanetary magnetic field, the field lines of which are wound into
helices by the Sun's rotation.

The region of expanding solar wind forms the
heliosphere, which is blown out like a bubble into
interstellar space and which is bounded at about 100-200 AU by the
heliopause. Initially the solar
wind flows at about 1.5 million km/h but it slows down as it ploughs into the galactic wind of neutral
hydrogen (and helium) atoms flowing from the galactic centre. At the heliopause, the pressures of the
solar and galactic winds are equal and they balance. Further beyond the heliopause is the
- a shock wave formed by the rapid deceleration of the galactic wind as it collides with the
heliosphere. The heliosphere is possibly drawn out into a comet-shaped tail as the Sun moves in its
own galactic orbit, ploughing through the galactic medium, although recent data (from IBEX) suggest
that there isn't much of a comet-shaped tail and the heliosphere is more bubble-shaped. [External
link to
The corona

The corona (lit. 'crown') is the outermost layer of the Sun's atmosphere. It is very diffuse and of
indefinite thickness (indeed it gradually thins as it fills much of the Solar System as the solar wind).
The corona is wispy, tenuous and whitish and not normally seen over the brilliance of the underlying
photosphere, but it is visible as a crown of spikes during a total eclipse.
Above: coronal mass ejection (CMEs) - plasma that erupts into space and leaves
the surface of the Sun. Click to enlarge. [Images courtesy of SOHO/NASA.]
Above: Solar prominences - plasma eruptions that are confined to coronal
magnetic field loops. Click to enlarge. The image on the left shows a double
prominence. {Images courtesy of SOHO/NASA.]
Coronal loops NASA
Above: a model of the coronal magnetic field showing loops and open flux lines. [Image courtesy of
A deuterium nucleus (deuteron) can then likewise fuse with a third proton to form a helium-3 nucleus:
Two helium-3 nuclei can then fuse to give a helium-4 nucleus and two protons (which can begin the
process again):
Overall the process may be represented in a single equation:
Overall mass is converted into energy! The first step in this reaction sequence, the fusion of two
protons, is very slow, so slow that the average lifetime of a proton inside the Sun (that is the average
length of time it exists in the Sun before colliding with another proton to form a deuteron) is 9 000 million
years - about the lifespan of the Sun! Indeed, this is just as well, for if it were faster than stars like the
Sun would be shorter lived and their planets would have less time to evolve life. This is the rate-limiting
step and creates a bottleneck in the process.

Why is it so difficult to join two protons in nuclear fusion?

Protons are positively charged, and like electric charges repel one-another. This means that it takes a
lot of energy to force two protons close enough together to fuse, since the closer they get to each other,
the stronger they repel each other! This electric barrier (or Coulomb barrier) is so hard to breach that at
temperatures less than 10 million K it barely occurs. Even at the high temperature's in the Sun's core,
this reaction would occur too slowly if it wasn't for a neat quantum trick -
quantum tunnelling. When
particles encounter a barrier, there is a small probability that they will effectively dematerialise and
materialise at the other side of the barrier - they tunnel through it! Thus, when two protons approach
close enough, even if they do not meet, one of them can dematerialise and then materialise in the
vicinity of the other, at the other side of the Coulomb barrier, and then both can react and fuse. The
Sun relies on this mysterious process of quantum tunnelling to keep it burning. They must get close
enough for tunnelling to work, so the high temperatures give the protons enough energy to come close
enough together to tunnel before being repelled.

The energy radiated away from the core is mostly in the form of photons, though some of it is radiated
as neutrinos. The photons travel through the relatively transparent
radiative zone. In this zone, there
is still enough opacity to make progress hard for the photons and they bounce from particle to particle,
making their way toward the convective zone in a slow random walk. It is estimated that it may take on
average 20 million years for a photon produced in the Sun's core to traverse the radiative zone! At the
convective zone, temperatures have fallen to about 1 million K and some electrons can recombine with
the nuclei and this gives them the ability to absorb photons - the plasma becomes opaque. In this region
temperature gradients build up and this drives heat transport by convection - this is the convective zone.
The diagram on the left (click to enlarge) is a simulation of
solar convection cells generated using the Pov-Ray object
collection 252452, 'thesun'. [External link to Pov-Ray

A series of these images have been animated and made into
an animated gif -
click here to view the animation.

The scale may not be right for solar convection (the cells are
much smaller) but it is a clever little graphic to illustrate the
general principle.