The major uncertainties involved in the Chandrasekhar mass models for Type Ia supernovae (SNe Ia) are related to the companion star of their accreting white dwarf progenitor (which determines the accretion rate and consequently the carbon ignition density) and the flame speed after the carbon ignition. We calculate explosive nucleosynthesis in relatively slow deflagrations with a variety of deflagration speeds and ignition densities to put new constraints on the above key quantities. The abundance of the Fegroup, in particular of neutron-rich species like 48 Ca, 50 Ti, 54 Cr, 54,58 Fe, and 58 Ni, is highly sensitive to the electron captures taking place in the central layers. The yields obtained from such a slow central deflagration, and from a fast deflagration or delayed detonation in the outer layers, are combined and put to comparison with solar isotopic abundances. To avoid excessively large ratios of 54 Cr/ 56 Fe and 50 Ti/ 56 Fe, the central density of the "average" white dwarf progenitor at ignition should be as low as < ∼ 2 × 10 9 g cm −3 . To avoid the overproduction of 58 Ni and 54 Fe, either the flame speed should not exceed a few % of the sound speed in the central low Y e layers, or the metallicity of the average progenitors has to be lower than solar. Such low central densities can be realized by a rapid accretion as fast aṡ M > ∼ 1 × 10 −7 M ⊙ yr −1 . In order to reproduce the solar abundance of 48 Ca, one also needs progenitor systems that undergo ignition at higher densities. Even the smallest laminar flame speeds after the low-density ignitions would not produce sufficient amount of this isotope. We also found that the total amount of 56 Ni, the Si-Ca/Fe ratio, and the abundance of some elements like Mn and Cr (originating from incomplete Si-burning), depend on the density of the deflagration-detonation transition in delayed detonations. Our nucleosynthesis results favor transition densities slightly below 2.2×10 7 g cm −3 .

We present an extensive study of the inception of supernova explosions by following the evolution of the cores of two massive stars (15 M ⊙ and 25 M ⊙ ) in multidimension. Our calculations begin at the onset of core collapse and stop several hundred milliseconds after the bounce, at which time successful explosions of the appropriate magnitude have been obtained. Similar to the classical delayed explosion mechanism of Wilson (1985), the explosion is powered by the heating of the envelope due to neutrinos emitted by the protoneutron star as it radiates the gravitational energy liberated by the collapse. However, as was shown by Herant, Benz & Colgate (1992), this heating generates strong convection outside the neutrinosphere, which we demonstrate to be critical to the explosion. By breaking a purely stratified hydrostatic equilibrium, convection moves the nascent supernova away from a delicate radiative equilibrium between neutrino emission and absorption. Thus, unlike what has been observed in one-dimensional calculations, explosions are rendered quite insensitive to the details of the physical input parameters such as neutrino cross-sections or nuclear

We present a new nucleosynthesis process that we denote as the nu p process, which occurs in supernovae (and possibly gamma-ray bursts) when strong neutrino fluxes create proton-rich ejecta. In this process, antineutrino absorptions in the proton-rich environment produce neutrons that are immediately captured by neutron-deficient nuclei. This allows for the nucleosynthesis of nuclei with mass numbers A>64, , making this process a possible candidate to explain the origin of the solar abundances of (92,94)Mo and (96,98)Ru. This process also offers a natural explanation for the large abundance of Sr seen in a hyper-metal-poor star.

Supernova simulations to date have assumed that during core collapse electron captures occur dominantly on free protons, while captures on heavy nuclei are Pauli-blocked and are ignored. We have calculated rates for electron capture on nuclei with mass numbers A = 65-112 for the temperatures and densities appropriate for core collapse. We find that these rates are large enough so that, in contrast to previous assumptions, electron capture on nuclei dominates over capture on free protons. This leads to significant changes in core collapse simulations. PACS numbers: 26.50.+x, 97.60.Bw, At the end of their lives, stars with masses exceeding roughly 10 M ⊙ reach a moment in their evolution when their iron core provides no further source of nuclear energy generation. At this time, they collapse and, if not too massive, bounce and explode in spectacular events known as type II or Ib/c supernovae. As the density, ρ, of the star's center increases, electrons become more degenerate and their chemical potential µ e grows (µ e ∼ ρ 1/3 ). For sufficiently high values of the chemical potential electrons are captured by nuclei producing neutrinos, which for densities 10 11 g cm −3 , freely escape from the star, removing energy and entropy from the core. Thus the entropy stays low during collapse ensuring that nuclei dominate in the composition over free protons and neutrons. During the presupernova stage, i.e. for core densities 10 10 g cm −3 and proton-to-nucleon ratios Y e 0.42, nuclei with A = 55-65 dominate. The relevant rates for weak-interaction processes (including β ± decay and electron and positron capture) were first estimated by Fuller, Fowler and Newman [1] (for nuclei with A < 60), considering that at such conditions allowed (Fermi and Gamow-Teller) transitions dominate. The rates have been recently improved based on modern data and state-of-the-art many-body models [2], considering nuclei with A = 45-65. (This rate set will be denoted LMP in the following.) Presupernova models utilizing these improved weak rates are presented in [3]. In collapse simulations, i.e. densities 10 10 g cm −3 , a much simpler description of electron capture on nuclei is used. Here the rates are estimated in the spirit of the independent particle model (IPM), assuming pure Gamow-Teller (GT) transitions and considering only single particle states for proton and neutron numbers be- During core collapse, temperatures and densities are high enough to ensure that nuclear statistical equilibrium (NSE) is achieved. This means that for sufficiently low entropies, the matter composition is dominated by the nuclei with the highest binding energy for a given Y e . Electron capture reduces Y e , driving the nuclear composition to more neutron rich and heavier nuclei, including those with N > 40, which dominate the matter composition for densities larger than a few 10 10 g cm −3 . As a consequence of the model applied in previous collapse simulations, electron capture on nuclei ceases at these densities and the capture is entirely due to free proto...