Stellar Evolution

When we look to the stars, we look to the past. Light always plays catch-up with time, and the stars we see are evidence of this. A wonderful anomaly, it means we can see the universe at varying stages of its life simply by looking a little further or closer to home.

The stars may all appear the same – save for slight variations in colour and luminosity – but that could not be further from true. We could look to one region of the sky and see a birth place of stars, and in another we may find the remains of stars which have long since died.

Understanding stellar evolution goes beyond the beauty of realizing that stars have real life spans like everything else in the cosmos. For one, all the elements (except hydrogen) that make up our human bodies are a direct result of stellar evolution. A star the size of a sun gives up its atoms to space in beautiful planetary nebulae, littering nearby planets with elemental winds rich in carbon and oxygen. A more massive store containing elements like calcium can die a violent supernova death, also spreading its atomic seed. So much of what we know and learn about our universe comes from studying stars, from their enigmatic births to their cataclysmic deaths.

Birth and Equilibrium

Stellar Nurseries and Protostars

Galaxies contain an interstellar medium which is made up of mainly gas and dust. The unique composition of this medium differs somewhat from galaxy to galaxy, but in a spiral galaxy such as our own, the interstellar medium usually consists of roughly 70% hydrogen. Helium makes up the bulk of the remainder. The medium is abundant with millions of particles per cubic centimetre. Interstellar medium also contains trace amounts of heavier elements ejected by stars as they age and die.

The medium isn’t thoroughly consistent. Molecular clouds are regions of the interstellar medium which are able to form molecules. Giant molecular clouds are vast, massive regions of gas less dense than the gas of stars. Though they contain 100 to 1 million solar masses they are quite cool; only a couple of degrees Kelvin. They are known to contain almost a hundred different molecules including ethyl alcohol, and carbon monoxide. Giant molecular clouds are potential stellar nurseries.

Stars form when gravity can force small regions of the cloud to contract to a higher density and a higher temperature state, though several factors have to be met for this to happen, and it cannot just occur spontaneously. Instead, passing shock waves from supernovae may cause the gas to contract, or collisions between giant molecular clouds might trigger star formation. Nearby stars can also ionize gas and drive it away to cooler and denser medium.

Once the cloud starts to contract, gravity draws all the atoms together in the centre. The atoms gain speed as they fall to the centre, but that isn’t sufficient to heat the gas enough to form stars. The motion of the falling atoms has to be highly randomized, colliding with one another at great velocities. The thermal energy from these random collisions increases the temperature of the gas. Gravitational energy is converted to thermal energy because the atoms speed up, collide and grow hot, all due to the downward movement of gravity exerting a force.

The clumps in giant molecular clouds can be called protostars when they are hot enough to radiate infrared but not yet hot enough for nuclear fusion to occur. Protostars have high density centres cocooned by low density clouds. The protostar is large and luminous as it contracts from the inside out, and mass moves inward from the outer parts of the cloud. It begins to rotate more rapidly and flattens the surrounding cloud into a disc. These protostar discs are where astronomers believe planets form. As the gas loses its momentum it continues to fall to the centre and the protostar grows. The protostar then drives surrounding gas and dust away with stellar winds and radiation as it becomes bigger and hotter, finally revealing itself from its nebulous cocoon.

The H-R Diagram and Main-Sequence Stars

The Hertzsprung-Russell diagram (H-R diagram) is a graph which plots stars by their temperature and surface area against luminosity. It sorts by size too. The stars are plotted with luminosity on the vertical axis and temperature on the horizontal axis.

There is a distinct band of stars running along a section of the graph scientists call the main sequence. The types of stars that lie on this band are the most abundant in the observable universe. A star’s main sequence phase begins when it starts fusing the hydrogen at its core into helium. The star is fuelled by these nuclear reactions.

Main sequence stars hit the balance between gravity which contracts the star, and internal thermal pressure which expands it. We say that such a star is in a state of hydrostatic equilibrium. There is also a special relationship between pressure and temperature, a built in thermostat to keep the star burning steadily. If a star’s nuclear reactions produce too much energy, the excess flows out of the star and makes its layers expand. The expanding layers lower the central temperature and density, and the nuclear reactions are slowed until they are stable again. Conversely, the star would contract if the nuclear reactions make too little energy, and the central temperature and density would rise until nuclear energy production is once again sufficient.

The main sequence phase lasts as long as the star is in hydrostatic equilibrium; generating thermal energy as it fuses hydrogen for fuel. Small stars burn through their hydrogen far more slowly than massive stars do, so a star’s lifespan and what happens after the main sequence depends on its mass.

Aging Stars

Giants and Dwarves

We have just seen that a star’s nuclear reactions generate energy which keeps the interior hot. This creates a high pressure environment which balances gravity and keeps the star from crushing under its weight. The reality is gravity always wins in the end. The star’s supply of hydrogen at the core eventually begins to dwindle, and this means the beginning of death for a star.

The main sequence lasts a few billion years for Sun-like stars (approximately 0.4 – 4 solar masses). Our own Sun’s main sequence phase is estimated to last about 10 billion years, and scientists believe we are half way through our star’s active life cycle. It is a long lifetime, but it does come to an end.

Imagine a fire burning without being stirred or adding wood from around the fire’s centre. The fire would die out and leave behind ash. Sun-like stars have little to no convection at their centres, and the same principle applies. An ash of inert helium nuclei collects at a star’s core as hydrogen fuses. The helium as is not hot enough to fuse into heavier elements and so the energy production falls. Hydrogen in the outer layer of the core can’t be ‘mixed’ because of the lack of convection. This gives gravity the edge, and the star’s core starts contracting under its weight.

An interesting thing then happens: the helium nuclei grow hotter because of the gravitational energy crushing the core. The gravitational energy is once again converted into thermal energy. The thermal energy now heats the outer shell of hydrogen which surrounds the core. Hydrogen in the shell will start fusing when the shell gets hot enough.

However, the hydrogen-fusing shell cannot stop the contractions of the helium core because the core has no energy source to counteract gravity. The gas in the core becomes extremely dense as it is squeezed together and compressed. In turn, the particles (atomic nuclei and free electrons) of the gas are also squeezed together. This poses a problem. Free moving electrons can only have a certain amount of energy, and can only exist at certain energy levels when bound to an atom. Electrons spin in one of two directions, and no two electrons of the same spin can occupy the same energy level. This means that as the electrons and atomic nuclei are squeezed together, the lower energy levels fill up quickly with electrons.

In this case, free electrons cannot slow down as gravity pushes against it and compresses the gas. The electron can only speed up to reach the higher energy levels. If the gas is so dense that the electrons are not free to change their energy, it is called electron degenerate matter. The electron degeneracy pressure is the pressure caused by this state of the electrons, and the pressure is enough to stop the core from collapse.

Finally, if the star’s core gets hot enough – 100 million Kelvin – helium will fuse into carbon. The star makes more energy than is needed from its hydrogen-fusing shells and helium-fusing cores. The excess energy radiates outward, expanding the layers of the star’s gas and forming a red giant.

Red dwarves are stars less massive than about 0.4 solar mass. They can live for an extremely long time because of the slow and steady rate at which they burn their hydrogen. Small stars like these are convective throughout their entire structure, meaning that as hydrogen is consumed, helium collects uniformly throughout the star. The star will never have a hydrogen fusing shell, and the star won’t ever become a red giant. When all the hydrogen is used up, an inert helium stellar remnant is left behind. Of course, we have not yet discovered such remains, nor could we. Estimates suggest that it could take 6 trillion years for a red dwarf to use up all its fuel. That means that every single red dwarf in the universe is still fusing hydrogen today because the universe is only 13.8 billion years old.

Planetary Nebulae

Planetary nebulae are made of the ionized gases shed by a dying star. When Sun-like stars swell into giants and age, they expel their outer atmosphere into space. Once the hot interior of the star is exposed, high stellar winds are also ejected. Ultraviolet radiation from the hot remains of the core excite the gas and cause it to glow. Their radii can reach 3 light years across. All the planetary nebulae we see are no more than around 10 000 years old. After this period, the shed gases become mixed into the interstellar medium. Scientists have records of about 1500 planetary nebulae, and because they are so short lived, we can conclude that they are quite common occurrences in medium mass stars.

White Dwarfs

White dwarves are the remains of medium mass stars which fused hydrogen and helium, but were not massive enough to fuse carbon. These stars have driven their outer layers away forming a planetary nebula, and then collapsed into a white dwarf.

The white dwarf star is made up of mainly carbon and oxygen ions among a sea of degenerate electrons. The degenerate electron matter supports the star against its own gravity because electrons are unable to pack into a tighter space. As you may have already determined, they are very dense objects. Their structure is unique. Though the degenerate electrons create the pressure needed to support the object, most of its mass comes from the carbon and oxygen ions. It generates no nuclear energy and contains no gas, so it is more accurate to call it a ‘compact object’ rather than a star.

The white dwarf’s compact core is very hot, but this heat slowly flows to the surface and escapes. Billions of years are needed for all the heat to radiate out of its small surface area. Given enough time, the object will turn cold and dark – a black dwarf. Our universe isn’t old enough to contain black dwarves; not enough time has passed for these stars to have radiated all their energy out into space.


Due to their size, more massive stars are able to get hot enough to fuse carbon and oxygen. This means that their fates are very different from medium and low mass stars. The red giant phase of a star’s life causes it to lose a lot of mass as its gasses are shed into space, but if the star is still over 4 solar masses when its carbon and oxygen core contracts, it can reach 600 million Kelvin. That is hot enough for carbon to fuse. The process continues much the same, fusing heavier and heavier elements until iron is made. The nuclear reactions of heavier elements happen relatively quickly. Hydrogen fusion can last a couple of million years, but the same star can fuse all its silicon in a day.

Iron cannot fuse, so the star builds up so much iron that its core alone can outweigh the sun. Eventually the core must collapse, and when it does, a supernova occurs. As the core rapidly collapses, a shockwave tears through the star, blowing apart with energy equivalent to 1028 megatons of TNT.

There are two types of supernovae explosions: type I and type II supernovae. Type II supernovae are the sort we have just examined; a supernova caused by the collapse of a massive star. Type Ia and type Ib supernovae both occur within binary star systems. A type Ia supernova is the result of a white dwarf gaining mass from its companion star. It eventually becomes so overwhelmed that a supernova which can be six times as bright as a type II event is triggered. Type Ib are less common. Astronomers believe these events happen when a massive star loses its outer layers to its companion, exposing the core. Otherwise, the star develops normally, until its iron core collapses in a supernova.

Contrary to what some may believe, supernovae happen rapidly, quietly and are remote and rare. We have recorded several supernovae throughout our universe, but only 1 or 2 happen within the Milky Way every century. When a star does go supernova, it is so energetic and luminous it can outshine its entire host galaxy. They are so bright that they can be seen during the day before they quickly fade back into obscurity.

Neutron Stars

Neutron stars are another type of compact object. They are the remains of a star which has gone supernova, and are made up of neutron degenerate matter.

The weight of a massive collapsing core cannot be supported by electron degeneracy pressure. As the core collapses, things are so hot and dense that atomic nuclei are broken apart by gamma rays. The electrons and protons that are freed during this process are forced together by the extreme density, and combine to form neutrons.

Neutron stars are much smaller, hotter and denser than white dwarves. It would take possibly billions of years for such an object to radiate all its heat from such a small surface area. Their magnetic fields are estimated to be possibly a trillion times that of the Sun. They are so dense that even a sugar cube sized of neutron star matter would weigh 100 million tons! Because such a great mass is focused in a small compact area, neutron stars spin very rapidly – about 10 to 100 times per second. For reasons not fully understood, some neutron stars emit timed bursts of energy as they rotate, and then they are known as pulsars.

Black Holes

Black holes: the most enigmatic end a massive star could endure. They are objects that turn the laws of physics upside, distorting space and time.

We know that with gravity, what goes up must come down. If you wanted to leave the surface of the Earth – escape its gravitational pull – you would have to move at 11km/s, like in a jet. This is the escape velocity you need to leave the surface of Earth. Black holes are so dense that no even light travelling at 300 000 km/s can escape.

Stellar black holes are left behind by the collapse of the most massive stars. For medium mass stars, electron degeneracy pressure stops the collapse of the core and they become white dwarves. For bigger stars, neutron degeneracy pressure halts the collapse of the star and it becomes a neutron star. For the most massive stars, no force remains to prevent the object from collapsing into a singularity: an object so compressed that it is said to have zero radius. A singularity has infinite density and infinite gravity.

Black holes and the regions of space near them are virtually undetectable. You can receive no information about such an object.  We can only observe them by the strange effects they have on the space around them.

Because we do not fully understand black holes, there are a lot of myths surrounding them. Firstly, they are not star-swallowing monstrosities! A black hole’s gravitational field will not affect objects unless you get extremely close to it. As a matter of fact, the Sun could be replaced by a black hole which has the same mass, and the planets would remain safe. But happens if you do get too close?

For a person falling toward a black hole, time may pass as normal. For outside observers, they would watch the doomed astronaut fall ever more slowly toward the black hole’s event horizon (we are unable to view any event which happens passed this boundary) until they appear to not be moving at all. Time and space stretch and distort near a black hole, so the astronaut may be stretched lengthwise by tidal forces, and simultaneously compressed laterally. On the other hand, the extreme environment would likely incinerate anything which got to close. These stellar corpses are so bizarre that some even think time travel may be possible with black holes!


It is safe to say that our understanding of stellar evolution and stellar remnants like black holes is not complete. There is something to be said about the genius and dedication of the scientists who study the stars which no one can ever dream of exploring – the very same stars from which we are made – and bringing the light of that knowledge to the general population.