Molecular cloud → core → protostar → star
An interactive, physics-first walkthrough of star birth — collapse criteria, timescales, the disks that surround newborn stars, how we detect those disks in the infrared, and how our own Solar System assembled. Hover anything ? for help.
The big picture
A star is the end of a long collapse. Diffuse interstellar gas gathers into a giant molecular cloud, fragments into dense cores, and — when gravity overwhelms pressure — falls inward to ignite hydrogen fusion. The whole story is a contest between two forces.
Self-gravity wants to compress the cloud. The more mass packed into a smaller volume, the stronger it gets — and the faster collapse accelerates once it begins.
Thermal pressure (set by temperature) and turbulence resist compression. Warm, low-density gas resists easily; cold, dense gas cannot. That is why stars form in the coldest clouds.
Astronomers classify a forming star by its infrared signature, which tracks how deeply it is still buried in gas and dust. Each stage has a characteristic duration — explore them below.
Durations are order-of-magnitude and vary with stellar mass: a massive O star runs through every phase faster, a low-mass M dwarf far more slowly.
The collapse criterion
In 1902 James Jeans asked a simple question: how big does a gas clump have to be before its own gravity beats its internal pressure? The answer is a critical mass and a critical length. Exceed them, and collapse is unstoppable.
Read it as a tug-of-war: warmer gas (larger T) raises the bar for collapse, while denser gas (larger ρ) lowers it. The companion Jeans length sets the critical size.
Drag the sliders to build a gas clump and watch the verdict.
The collapse clock
Once a core is unstable, how fast does it fall in? If pressure is negligible, the answer depends on density alone. Denser gas collapses faster — which is why collapse accelerates as it proceeds (runaway collapse).
Strikingly, mass and radius cancel out — only density matters. A key consequence: because density rises during collapse, the centre runs away first, seeding a protostar while the outer envelope is still falling.
Pressure slows real collapse somewhat, so t_ff is a lower bound — but it is the right order of magnitude for the prestellar and Class 0 phases.
Diffuse cloud → core → protostar: each rise in density shortens the clock by orders of magnitude.
What surrounds a newborn star
A collapsing cloud always has a little spin. As it shrinks, conservation of angular momentum forces it to spin faster, and gas settles into a flattened, rotating protoplanetary disk around the protostar. These disks are the birthplaces of planets.
~99% of the mass — mostly H₂ and He. Drives accretion onto the star and forms the envelopes of giant planets before it disperses.
~1% by mass: silicate and carbon grains. Sticks, grows, and settles to the midplane — the seed material for rocky planets and planet cores.
Beyond the snow line, volatiles (H₂O, CO, CO₂) freeze onto grains, boosting solid mass and feeding giant-planet cores.
Hot (>1000 K) near the star, dropping to tens of K at the edge. This gradient is exactly what makes disks glow in the infrared — see the next section.
A passively heated disk follows a power law in radius. Inner regions are hot enough to sublimate dust; outer regions are frozen.
How we find the disks
A bare star emits like a single blackbody — its spectrum peaks at one wavelength set by its surface temperature. But warm dust in a surrounding disk re-radiates absorbed starlight at longer, cooler wavelengths. The result is more infrared light than the star alone can explain: an infrared excess. That bump is the disk's fingerprint.
Toggle the disk and reshape it. The axes are log–log: wavelength across, energy emitted up.
This is the principle behind IRAS, Spitzer, WISE and JWST disk surveys: photometry at several IR bands reveals the excess over the stellar photosphere.
Worked examples
From buried protostars to the Sun itself. Click any object to open its full profile: physical parameters, disk state, infrared signature, and an estimate of how long it lives.
Lifetime estimates use the main-sequence scaling t ≈ 10¹⁰ yr × (M/M☉)⁻²·⁵, derived from L ∝ M³·⁵ and a fixed fuel fraction. Hover any number in a profile for its meaning.
Our own origin
Everything above culminates here. About 4.57 billion years ago, a fragment of a molecular cloud collapsed, spun up into a disk, and built the Sun and planets. The same physics — collapse, disks, a snow line, temperature-sorted material — explains the layout we see today.
A dense core in a molecular cloud — perhaps triggered by a nearby supernova shock — exceeds its Jeans mass and begins to fall in under its own gravity.
Tiny initial rotation is amplified as the cloud shrinks. Gas cannot fall straight to the centre, so it flattens into the rotating solar nebula — a disk feeding a growing protosun.
The dense centre heats until hydrogen fusion begins. A fierce T Tauri wind later blows away the remaining gas.
Material condenses by temperature. Close in, only metals and silicates survive; past the snow line, ices condense too — far more solid mass is available out there.
Dust sticks into pebbles, pebbles into kilometre-sized planetesimals, which accrete into protoplanets through gravity.
Rocky terrestrial planets inside the snow line; ice-and-gas giants beyond it, where large cores could capture nebular gas before it dispersed. Leftovers became asteroids and comets.
The snow line (~2.7 AU in the early Solar System, where T ≈ 150–170 K) split the disk into a rock-only inner zone and an ice-rich outer zone.
Quick reference