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Every star, ancient gas cloud, and distant galaxy in the observable universe consists of roughly 75% hydrogen and 25% helium-4 by mass. Stars fuse hydrogen into helium, but stellar fusion has only altered this ratio by about two percent over 13.8 billion years.
Where did this universal 3-to-1 mass ratio come from if stars didn't make it?

The First Second and Freeze-Out

In 1948, physicists Ralph Alpher and George Gamow proposed that all atomic nuclei formed during the scorching first minutes following the Big Bang, a process called Big Bang Nucleosynthesis (BBN).
When the universe was under one second old, temperatures exceeded 1010Β K10^{10}\text{ K} (kBT>0.8Β MeVk_B T > 0.8\text{ MeV}), allowing weak force reactions to rapidly interconvert protons (pp) and neutrons (nn) through collisions with electrons and neutrinos:
n+Ξ½e↔p+eβˆ’andn+e+↔p+Ξ½Λ‰en + \nu_e \leftrightarrow p + e^- \quad \text{and} \quad n + e^+ \leftrightarrow p + \bar{\nu}_e
Because a neutron is slightly heavier than a proton by Ξ”mc2=1.293Β MeV\Delta m c^2 = 1.293\text{ MeV}, the equilibrium ratio of neutrons to protons is governed by the Boltzmann factor:
(np)eq=eβˆ’Ξ”mc2kBT\left(\frac{n}{p}\right)_{\text{eq}} = e^{-\frac{\Delta m c^2}{k_B T}}
As the universe expanded and cooled below 0.8Β MeV0.8\text{ MeV}, the rate of weak interactions dropped below the expansion rate of space, causing the ratio to freeze out at approximately n/pβ‰ˆ1/6n/p \approx 1/6.
Once free neutrons are cut off from rapid conversion back from protons, what stops them from simply decaying away?

The Deuterium Bottleneck

Free neutrons are unstable and undergo beta decay (nβ†’p+eβˆ’+Ξ½Λ‰en \rightarrow p + e^- + \bar{\nu}_e) with a mean lifetime of Ο„nβ‰ˆ880Β seconds\tau_n \approx 880\text{ seconds} (about 14.6 minutes).