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A star is not a fixed object. Its properties — luminosity, temperature, radius, composition — change over time, sometimes dramatically. When its nuclear fuel runs low, a star cannot simply keep shining; it must rearrange itself to find a new source of energy, and in doing so it moves across the Hertzsprung–Russell diagram.
This lesson traces the life story of a star from birth, through main-sequence burning, to death. The details depend strongly on the star's initial mass: low-mass stars (like the Sun) end as white dwarfs, while high-mass stars die spectacularly as supernovae and leave behind neutron stars or black holes. Along the way, the star's position on the HR diagram changes in characteristic ways — the evolutionary track.
This is specification content 5.5.2, and it is one of the most conceptually important topics in OCR Module 5.5. You will be asked to describe the main stages of stellar evolution, explain what physical processes drive each transition, and relate the stages to positions on the HR diagram.
Stars form from collapsing clouds of interstellar gas and dust — giant molecular clouds, consisting mostly of hydrogen with trace amounts of helium and heavier elements. A dense clump in such a cloud, triggered by a passing shockwave or simply by its own gravity, begins to collapse under its own weight. As it collapses:
When the core temperature reaches about 10⁷ K, hydrogen fusion via the proton–proton chain (or the CNO cycle in more massive stars) ignites. The new star releases energy fast enough to counterbalance gravitational collapse, and it settles into hydrostatic equilibrium: outward radiation pressure and gas pressure balance inward gravity. A main-sequence star is born.
OCR does not ask for the details of star formation in Module 5.5, so we shall not dwell on it further. The key point is that every star begins its main-sequence life with a mass set by its protostellar cloud, and that mass determines everything that follows.
Once on the main sequence, a star fuses hydrogen to helium in its core. This is by far the longest phase of its life. The main-sequence lifetime depends on mass:
t_MS ∝ M / L ∝ M / M^3.5 = M^(-2.5)
A more massive star has more fuel, but it burns that fuel much faster. The net result is that a 10-M_☉ star lives only about 10⁷ years on the main sequence, while a 0.5-M_☉ red dwarf would last 10¹² years — longer than the current age of the universe.
| Mass (M_☉) | Main-sequence lifetime (years) |
|---|---|
| 50 | approx 5 × 10⁶ |
| 10 | approx 2 × 10⁷ |
| 1 (Sun) | approx 10¹⁰ |
| 0.5 | approx 10¹¹ |
| 0.1 | approx 10¹² |
The Sun is roughly half-way through its main-sequence life: approx 4.6 billion years so far, with another 5 billion to go. The slowest-evolving stars are the ones that have the longest future ahead of them.
When a low-mass star exhausts the hydrogen in its core, energy generation stops in the very centre. Without the outward pressure of fusion, the core begins to contract under gravity. Contraction heats the surroundings; in a thin shell just outside the dead core, temperatures rise high enough to ignite hydrogen-shell burning.
This has a strange consequence. The shell burns vigorously, releasing more energy than the old core did, but the energy has to be transported out through the stellar envelope. The envelope expands to adjust. As it expands, it cools — and the star becomes a red giant.
graph LR
MS[Main Sequence<br/>H burning in core] --> SGB[Subgiant Branch<br/>H shell burning]
SGB --> RGB[Red Giant Branch<br/>envelope expands]
RGB --> HB[Helium flash<br/>core He burning]
HB --> AGB[Asymptotic Giant<br/>double shell burning]
AGB --> PN[Planetary Nebula<br/>envelope ejected]
PN --> WD[White Dwarf<br/>cooling C/O core]
On the HR diagram, the star moves from the main sequence upward and to the right — it becomes much more luminous (because of shell burning releasing a great deal of power) and cooler at the surface (because the envelope has expanded). This is the red giant branch.
Eventually, the contracting helium core becomes hot enough (approx 10⁸ K) to ignite helium fusion via the triple-alpha process (3 ⁴He → ¹²C). In lower-mass stars this ignition happens in a dramatic event called the helium flash, because the core is partially electron-degenerate and temperature runs away briefly before thermal pressure restores order. In higher-mass stars the transition is smoother. Either way, the star settles into a new equilibrium burning helium in its core and hydrogen in a shell above.
After core helium is exhausted, the star climbs the asymptotic giant branch (AGB), burning helium and hydrogen in alternating shells around an inert carbon–oxygen core. The outer layers become unstable, pulsate, and gradually drift away from the star on stellar winds, eventually ejecting a planetary nebula — a glowing shell of gas illuminated by ultraviolet light from the hot exposed core. (The name is a historical misnomer: they have nothing to do with planets. Herschel called them "planetary" because their small greenish discs looked superficially like Uranus through early telescopes.)
What remains at the centre is a hot, dense, Earth-sized white dwarf made mostly of carbon and oxygen. It contains no fusion; its luminosity comes from residual heat, which it radiates away over billions of years, slowly fading in the lower-left corner of the HR diagram.
The white dwarf cannot be arbitrarily massive. Its support against gravitational collapse comes from electron degeneracy pressure — a quantum-mechanical effect in which electrons, forbidden by the Pauli exclusion principle from occupying the same state, resist being compressed beyond a certain limit. But degeneracy pressure has an upper bound: if the mass exceeds
M_Ch ≈ 1.4 M_☉
then electron degeneracy cannot support the star and it collapses further. This is the Chandrasekhar limit, derived by Subrahmanyan Chandrasekhar in 1931 (for which he received the 1983 Nobel Prize in Physics). Any stellar remnant exceeding 1.4 M_☉ cannot be a stable white dwarf.
For stars with initial main-sequence mass less than about 8 M_☉, mass loss during the red-giant and AGB phases reduces the final core to well below 1.4 M_☉. These stars therefore die as white dwarfs, and their fate is sealed.
Massive stars live fast and die young. Their main-sequence lifetimes are measured in tens of millions of years rather than billions. When their core hydrogen is exhausted, they follow a similar pattern to low-mass stars — shell burning, envelope expansion — but on a grander scale. They become red supergiants, with radii of hundreds to thousands of solar radii.
Unlike low-mass stars, they can ignite successive shells of heavier elements as their cores contract and heat up further:
graph TD
A[Hydrogen burning] --> B[Helium burning<br/>→ C, O]
B --> C[Carbon burning<br/>→ Ne, Mg]
C --> D[Neon burning<br/>→ O, Mg]
D --> E[Oxygen burning<br/>→ Si, S]
E --> F[Silicon burning<br/>→ Fe]
F --> G[Core collapse<br/>→ Supernova]
Each successive stage happens faster than the one before. Carbon burning takes about a thousand years for a 25-M_☉ star; neon burning lasts a few years; oxygen burning a few months; silicon burning only a day. By the end, the star has an onion-like structure with shells of successively heavier elements burning around an iron core.
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