Stellar Evolution: Stages in the Life Cycles of Stars
By Dr. Thomas Swan
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The Physical Features of Stars
The Physical Features of Stars
Stars are luminous spheres of fusing gas that are between 10 and 200,000 times wider than the Earth. The Sun is the nearest star to Earth and is 109 times wider. For an object to qualify as a star, it must be large enough for gravitational forces to have created enough pressure in its core to trigger nuclear fusion.
The surface temperature of the Sun is 5,500 °C, with a core temperature as high as 15 million °C. There are many other kinds of stars in the universe, with surface temperatures ranging from 2,500 °C to 50,000 °C. Stars are predominantly composed of hydrogen (73%) and helium (25%) gases, with traces of heavier elements such as oxygen, carbon, neon and iron.
Some smaller stars have lived since the earliest era of the universe, showing no signs of dying after more than 13 billion years. Others (the biggest and hottest) live only a few million years before burning through their fuel. These stars can have 300 times the mass (weight) of the Sun and be 7 million times as luminous. Conversely, the smallest stars can be 1/10th of the mass and have 1/10,000th of the luminosity.
Without stars, we would simply not exist. These cosmic behemoths use nuclear fusion to produce the elements that planetary systems and living beings are made of (carbon, oxygen, iron, etc.). The next sections will describe this process by outlining the different stages in the life cycles of stars.
The Birth of Stars
The Birth of Stars
Stars are born when nebulous clouds of hydrogen and helium gas coalesce under the force of gravity (see below). Sometimes a shock wave from a nearby supernova is required to produce areas of high density in the cloud.
As these dense pockets of gas accumulate material, gravity pulls them inward. The resulting compression heats the material, generating an outward pressure that slows the collapse. Once these opposing forces balance, the object reaches hydrostatic equilibrium; a gas pocket approaching this state is called a protostar.
A region of the Carina Nebula called Mystic Mountain in which stars are forming. ESA/Hubble via Wikimedia Commons.
Contraction comes to a complete stop when the core of the protostar becomes hot enough for hydrogen to fuse together to make helium in a process called nuclear fusion. At this point, the protostar becomes a main sequence star.
Star formation often occurs in gaseous nebulae in which the density of the nebula is great enough for hydrogen atoms to chemically bond to form molecular hydrogen. These nebulae or "clouds" are often called stellar nurseries because they contain enough material to produce thousands of stars, leading to the formation of star clusters (see below).
The Life of Stars
The Life of Stars
Once a star has formed and is on the main sequence, it is fusing hydrogen in its core. Hydrogen is the simplest form of atom, with one positively charged particle called a proton orbited by a negatively charged electron, although the electron is lost due to the intense heat of the star. This means that the hydrogen is ionized.
The fusion reaction fuelling the universe, in which four hydrogen nuclei (protons) fuse into one helium nucleus (He). Source.
The stellar furnace causes the ionized hydrogen (i.e., protons or 1H) to slam into each other. At core temperatures above 4 million °C, they fuse together to form helium (4He), releasing their stored energy in a process called nuclear fusion (see image).
During fusion, some of the protons are converted into neutral particles called neutrons in a process called radioactive decay (beta decay), which emits a small positron particle to carry away the proton's positive charge. The energy released in fusion heats the star further, causing more protons to fuse in a self-sustaining reaction.
Nuclear fusion continues in this sustainable way for between a few million and several billion years, until the star runs out of fuel (which may be longer than the age of the universe—13.8 billion years).
Contrary to expectations, the smallest stars (red dwarfs) live the longest. Despite having more hydrogen fuel, large stars (blue giants, supergiants, and hypergiants) burn through it quicker because the stellar core is hotter and under greater pressure from the weight of its outer layers. Smaller stars also make more efficient use of their fuel, utilising convective heat transport to circulate the fuel throughout the star's volume.
These are binary red dwarf stars (Gliese 623) that are 26 light years from Earth. The smaller star is only 8% of the Sun's diameter. Source.
If the star is large enough and hot enough (core temperature above 100 million °C), the helium produced in nuclear fusion reactions will also be fused together to form heavier elements such as carbon, oxygen, neon, and finally iron.
Elements heavier than iron, such as lead, gold, and uranium may be formed by the rapid absorption of neutrons, which then beta decay into protons. This is called the r-process for "rapid neutron capture," which is believed to occur primarily in neutron star mergers, although some contribution may come from supernovae.
The Death of Stars
The Death of Stars
Stars eventually run out of fuel. This first occurs in the stellar core because this is the hottest and densest region. Without the heat generated by fusion, the core begins a gravitational collapse, which consequently creates more heat and pressure, triggering fusion in the outer layers of the star where hydrogen fuel still remains.
These outer layers then expand to dissipate the heat, becoming larger and more luminous, but cooler on the surface due to the expansion (i.e., becoming redder). This is called the red giant phase of a star's evolution. Stars smaller than about 0.3 solar masses skip the red giant phase because they cannot become hot enough.
The contraction of the stellar core eventually becomes so extreme that the outer layers of the star are thrown off altogether, forming a planetary nebula (see below; although there are no actual planets).
The core stops contracting once the density reaches a point where electrons are prevented from moving any closer together. This is called "electron degeneracy" and it is determined by a physical law called Pauli's Exclusion Principle. The exposed core is called a white dwarf, which gradually cools to become a black dwarf.
Stars of more than ~10 solar masses typically undergo a more violent expulsion of the outer layers called a supernova (see below). Inside these larger stars, the gravitational collapse generates densities high enough for electrons to be captured by protons, converting them into neutrons and emitting a burst of neutrinos. The energy released by the gravitational collapse powers the explosion.
The super-dense neutron core that is left behind is called a neutron star, which is prevented from collapsing further by neutron degeneracy in the same way white dwarfs are sustained by electron degeneracy.
Massive stars above about 20 solar masses leave behind cores/remnants that are too dense and heavy for neutron degeneracy to prevent them from collapsing further into a gravitational singularity. These stars end their lives as black holes.
The expulsion of a star's matter returns it to the cosmos, providing fuel for the creation of new stars. As larger stars contain heavier elements (e.g. carbon, oxygen, and iron), supernovae seed the universe with the building blocks for Earth-like planets and living beings.
Tracking the Evolution of Stars
Tracking the Evolution of Stars
As stars progress through life, their size, luminosity, and radial temperature change according to predictable natural processes that can be plotted on Hertzsprung-Russell diagrams (plots of luminosity against surface temperature). The following (more technical) sections describe this evolution.
How a Protostar Becomes a Star
How a Protostar Becomes a Star
Prior to igniting fusion and becoming a main sequence star, a contracting protostar will reach hydrostatic equilibrium at a surface temperature of ~3,500 °C. This relatively luminous state is followed by evolutionary stages called the "Hayashi track" and the "Henyey track," which are illustrated in the diagram below.
This is a Hertzsprung-Russell diagram that tracks the early evolution of the Sun from protostar to main sequence star. The evolution of heavier and lighter stars are compared.
As the protostar joins the Hayashi track, its luminosity begins to decrease. This is because it is losing energy—the energy it radiates is not being replaced by nuclear fusion yet, which only occurs when it joins the main sequence. This loss of energy causes the protostar to slowly shrink under gravity; as its radius decreases, so does its luminosity. The rate of contraction is slow because it is in hydrostatic equilibrium. The rate is therefore determined by how quickly energy can escape, which is limited by the protostar's high opacity.
The opacity is high because the protostar's cool atmosphere contains easily ionized metals (sodium, potassium), inherited from the stellar nursery. Their ionization supplies free electrons that bind with hydrogen to produce H– ions, which are opaque and trap heat, maintaining a near-constant surface temperature as the protostar contracts. As the contraction heats the core, the opacity makes radiative heat transport inefficient, so the protostar becomes fully convective—the defining property of the Hayashi track—with hotter material rising toward the surface.
If the protostar has accrued less than 0.5 solar masses, it will remain convective and stay on the Hayashi track for more than 100 million years before igniting hydrogen fusion and becoming a main sequence star. If a protostar has less than 0.08 solar masses, it will never reach the temperature required for nuclear fusion (i.e., never join the main sequence) and will end life as a brown dwarf. Protostars heavier than 0.5 solar masses develop a radiative core after a few million years (less for higher masses) and then join the Henyey track.
On the Henyey track, the protostar's core has become hot and dense enough that H⁻ ions dissociate, causing the opacity to decrease and radiative heat transfer to become dominant. This evolution allows the surface temperature to increase (as in the H-R plot), which should increase the luminosity. However, the protostar is simultaneously shrinking, and the decreasing radius compensates for the rising surface temperature, so the luminosity stays roughly constant.
The protostar continues to contract under gravity until the core temperature reaches several million °C (the exact value depends on the star's mass) and hydrogen fusion ignites. More massive stars complete this phase faster, reaching the main sequence in as little as 0.1 million years, while a star like the Sun takes tens of millions of years.
How a Main Sequence Star Evolves
How a Main Sequence Star Evolves
Once hydrogen fusion begins, stars enter the main sequence at a position dependent on their mass. Larger stars enter at the top left of the Hertzsprung-Russell diagram (see above), while smaller red dwarfs enter at bottom right.
The largest stars remain on the main sequence for only a few million years, while the smallest stars remain for perhaps trillions. The Sun will remain for 10 billion years (its current age is 4.5 billion).
When stars between 0.5 and 10 solar masses (e.g., the Sun) begin to run out of fuel, they leave the main sequence, becoming red giants. Stars larger than about 10 solar masses expand into red supergiants and undergo successive stages of nuclear fusion before their iron cores collapse in supernova explosions.
Indeed, after leaving the main sequence, stars larger than the Sun eventually heat their cores enough to fuse helium, and in the most massive stars, successive stages of fusion produce progressively heavier elements up to iron. These form rings like a tree, with hydrogen being the outer ring, then helium, then increasingly heavier elements toward the core (up to iron).
These are fusion shells within a massive star. At the center is iron (Fe). Shells are not to scale. Rursus via Wikimedia Commons.
This Hertzsprung-Russell diagram tracks the evolution of the Sun after it leaves the main sequence.
Red giants become particularly luminous due to their increased size and heat generation following the gravitational contraction of their cores. However, as their surface area is now much larger, their surface temperature decreases substantially (i.e., becoming redder). They move towards the top right of the Hertzsprung-Russell diagram (see image).
In stars between 0.5 and 2 solar masses, the contracting helium core becomes electron-degenerate—like in white dwarfs—so it cannot expand and cool. When the core temperature finally reaches ~100 million °C, helium fusion ignites in a runaway burst called a helium flash. The energy released lifts the degeneracy, the core expands, and the star briefly reverses its red giant phase. After the core helium is exhausted, it re-expands into an even more luminous giant.
In stars above 2 solar masses, the core is already hot enough to ignite helium before degeneracy sets in, so helium fusion begins smoothly with no flash.
How a Star Dies
How a Star Dies
Once all possible fuel is burnt, the core contracts under gravity, heating up in the process, until it is halted by degeneracy pressure (electron or neutron) or collapses to a singularity. Cores of less than 1.4 solar masses become white dwarfs (see below), which slowly cool to become black dwarfs. When the Sun becomes a white dwarf, it will be compressed to the size of the Earth and have about 50% of its original mass.
As alluded to earlier, cores heavier than 1.4 solar masses (Chandrasekhar limit) are compressed into 10-20 km wide neutron stars, and cores greater than approximately 2.3 solar masses (Tolman-Oppenheimer-Volkoff limit) become black holes.
It is possible for these objects to subsequently exceed these limits. A white dwarf that exceeds the Chandrasekhar limit through accretion will be destroyed in a supernova, while a neutron star that accretes enough matter to exceed the TOV limit will collapse into a black hole.
In most cases, the outer layers are expelled, forming a planetary nebula if the remnant is a white dwarf, or a supernova remnant if it is a neutron star or black hole.
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