At the end of its life, a massive star has a core composed mostly of iron-group nuclei. The core is surrounded by an onion-skin structure of shells dominated by successively lighter elements. Nuclear fusion is still ongoing in the shells, but the iron core is inert. The electrons in the core are relativistic and degenerate. They provide the lion’s share of the pressure support stabilizing the core against gravitational collapse. In this, the iron core is very similar to a white dwarf star, the end product of low-mass stellar evolution. Once the iron core exceeds its maximum mass (the so-called effective Chandrasekhar mass of $\sim$$1.5-2$ solar masses \[$M_\odot$\]), gravitational instability sets in. Within a few tenths of a second, the inner core collapses from a central density of $\sim$$10^{10}\,\mathrm{g\,cm}^{-3}$ to a density comparable to that in an atomic nucleus ($\gtrsim 2.7\times
  10^{14}\,\mathrm{g\,cm}^{-3}$). There, the repulsive part of the nuclear force causes a stiffening of the equation of state (EOS; the pressure–density relationship). The inner core first overshoots nuclear density, then rebounds (“bounces”) into the still collapsing outer core. The inner core then stabilizes and forms the inner regions of the newborn protoneutron star. The [hydrodynamic]{} supernova shock is created at 