Supernova Neutrinos

Neutrino physics: Theory and experiment (SS2021) Supernova neutrinos Teresa MarrodánUndagoitia Max-Planck-Institut fürKernphysik, Saupfercheckweg 1, 69117 Heidelberg, Germany E-mail: [email protected] Contents 1 Lecture 4: Supernova neutrinos2 1.1 Evolution of stars...............................2 1.2 Evolution of a heavy star...........................3 1.3 Core-collapse supernova...........................3 1.4 Neutrino emission in a supernova explosion.................6 1.5 Supernova 1987a...............................7 1.6 Neutrino detection.............................. 10 1.7 Summary................................... 11 1 LECTURE 4: SUPERNOVA NEUTRINOS 1. Lecture 4: Supernova neutrinos This lecture is about the production of neutrinos in supernovae and their detection on Earth. So far only one such event could be detected, the neutrinos from Supernova 1987a. We will discuss this measurement and the prospects for existing detectors. 1.1. Evolution of stars Most stars (like our Sun) have a mass M < Mcritical ' 8M . Their general characteristics are: • Lifetime of ∼ 1010 y • Temperature of ∼ 107 K (in the Sun) • Luminosity is mostly in photons (98% photo-luminosity and 2% neutrino luminosity) • Star end its life in a white dwarf More massive stars in which M > Mcritical have significantly different characteristics: • Shorter lifetime of ∼ 107 y 8 • Central temperature of & 10 K • Radiate mostly in neutrinos • Its life ends with a dramatic collapse of its core Figure1 shows an illustration of the evolution of stars in the classes described above. Figure 1. Life cycle of a star. Figure credit: NASA. 2 1 LECTURE 4: SUPERNOVA NEUTRINOS The probability for a core-collapse supernova in our Milky Way is about 3 in 100 years, so it is a rare event. 1.2. Evolution of a heavy star The burning stages of a heavy star are described in the following: • Hydrogen burning: In the first stage hydrogen is burned to helium as we saw in the solar neutrino lecture. This is the longest burning phase ∼ 7 · 106 y and the typical temperature is T = 0:02 · 109 K (see figure2). • When all hydrogen is burned, the energy production cannot withstand gravitation and the star contracts. Consequently, the pressure and the temperature increase. • The increased temperature allows for another burning stage. • Helium burning: helium is then burned to heavier elements: He ! 12C; 16O; 22Ne (1) • The typical time in this stage is 5 · 105 y and a temperature of T = 0:2 · 109 K. Once all He is burned, the star contracts and pressure and temperature increase. • Carbon burning: 12C ! 20Ne; 24Mg and similarly 22Ne and 16O burning. The temperatures in carbon, neon and oxygen burning range from 0:8 · 109 to about 1:8 · 109 K. • Silicon burning: this is the final stage in which silicon is burned to Fe, Ni, Cr. The burning time is about 1 day and T = 3:5·109 K. No further energy gain through fusion is possible as the highest binding energy (8 MeV/nucleon) is for iron (Fe). While for small stars only hydrogen fusion takes place, for large stars (M < 8M ) fusion up to iron is possible. Burning regions appear in an onion-like structure as shown in figure2. 1.3. Core-collapse supernova A core-collapse follows when a heavy star has consumed all available iron. The pre- condition is that the mass of the iron core shoud be larger than the Chandrasekhar mass: 2 Mcore > MCh = 5:7 · Ye · M (2) where Ye represent the number of electrons per nucleon and M a solar mass. The currently standard value of the Chandrasekhar limit is ∼ 1:4 M . As an example, for a 15 M : • MCh ∼ 1:5M 3 1 LECTURE 4: SUPERNOVA NEUTRINOS Figure 2. Evolution of a heavy star. Image credit: J. Hester & others. 9 • Tcentral ∼ 8 · 10 K • ρ ∼ 3:7 · 109 g/cm3 ∗ • Ye = 0:42 ∗ The Fermi energy of electrons is (4 − 8) MeV which are high enough to allow the following reactions: ( − e + p ! n + νe − Z Z−1 (3) e + A ! A + νe Once all iron in the core has been consumed, no burning process can compensate the gravitational force and the core collapses: 4 1 LECTURE 4: SUPERNOVA NEUTRINOS • The number of electrons is strongly reduced due to the electron capture on proton and photodissociation of iron reactions, for instance. This corresponds to stage B in figure3. D Shock propagation and e burst e e e A Progenitor E Shock stagnation and heating M M free n, p explosion 100 km H M M core He heating e X e+ p ~0.1 s A R p n S e n R cooling ~10 ms _ e,µ, e e,µ, shockgainradius R R ~200 km PNS G S gain layer (R~20 km) C Bounce and shock formation B Collapse of core XA ~1000 km nucleons e ~100 ms electron capture 12 3 and trapping at ~10 g/cm c nuclear photodisintegration matter ~1014 g/cm3 e core ~1.4 M c RS ~10 km Figure 3. Scheme of a supernova explosion. Figure from F. Kitaura. • The neutrinos produced in these reactions carry out energy from the core. • As the electrons were balancing the gravitational force, the core collapses quickly and the density rapidly increases. • For core densities of ρ ∼ 1012 g/cm3, ν-diffusion becomes larger than the collapse time and the neutrinos are trapped. The coherent neutrino-nucleus scattering process plays an important role at this stage. • The core does not cool down through neutrino loss anymore, the neutrinos thermalize and are also trapped by neutrons • Neutrino reactions taking place in the dense core matter are: 5 1 LECTURE 4: SUPERNOVA NEUTRINOS ν + A ! ν + A ν + (p; n) ! ν + (p; n) ν + e− ! ν + e− ν + ν ! e+ + e− (4) − + νe + n ! e + p νe + p ! n + e • The core finally reaches densities of ρ > 1014 g/cm3 (C in figure3). Infalling material bounces back and forms an outward-propagating shock front. • Matter collides with in-falling matter and produces a pressure wave which travels from the iron core to the exterior (D). The outer matter layers are blasted away while the inner remnant collapses to a neutron star or a black hole (E). The energy release in a supernova can be calculated by: GM 2 Gm2 ∆E = − − − (5) R star r neutron star the second term dominates due to the large difference in radius. The energy release can therefore be also written as 10 km m 2 ∆E = 5:2 · 1053 erg − · − neutron (6) rneutron 1:4 M where 1 erg = 10−7 J = 624 GeV. 1.4. Neutrino emission in a supernova explosion In collapsing stars, neutrinos and anti-neutrinos are produced by a variety of processes. Figure4 shows a prediction of the time evolution of the neutrino emission. Expected time dependence of the neutrino emission: • Sharp initial rise in luminosity of νe from rapid electron capture (B in figure3) 51 • About 5·10 erg energy are released in ∼ 15 ms by the νe in this spike. The energy of the neutrinos reach up to ∼ 12 MeV • From t ∼ 0:36 s, the intensity of all neutrino flavours rise. These neutrinos are emitted cooling down the produced proto-neutron star • The luminosities of νe and νe are very similar • In contrast, the luminosities of all other neutrinos are lower • Neutrino production processes: N + N ! N + N + ν + ν (7) + − e + e ! νe + νe (8) γ ! ν + ν (9) 6 1 LECTURE 4: SUPERNOVA NEUTRINOS Figure 4. Time evolution of the neutrino emission in a supernova explosion. The mean energy of each neutrino type is shown in the lower panels. Figure from [1]. This phase take a few seconds. This short time is important experimentally as it allows to reject largely backgrounds by selecting a short time window. 99% of the energy released by a supernova explosion is in the form of neutrinos. The mean energy of the neutrinos (see bottom of figure4) can be calculated considering the SN density, temperature, radius where they are emitted and the cross section. • < Eνe > ∼ 10 MeV • < Eν > ∼ 15 MeV for νµ and ντ as the production region is smaller compared to the one of νe 1.5. Supernova 1987a On February 23 rd, 1987, a supernova with the given name SN1987A exploded. It was located in the large Magellanian cloud, a companion galaxy of the Milky Way at ∼ 50 kpc (1:5 · 1018 km) from the Earth. 7 1 LECTURE 4: SUPERNOVA NEUTRINOS SN1987A was the brightest SN since the one that Kepler detected in 1604 and it could be detected in all wavelengths and also in neutrinos! • Progenitor star: a 20M blue supergigant • Large mass ejection in the explosion (see the image from Hubble in figure5) • Total explosive energy amounted to 1:4 · 1051 erg Figure 5. Picture of the SN1987a rest by Hubble telescope. Detection of SN1987A: • X-rays: the light curve of the explosion was measured by the ROSAT satellite • γ-rays: observed for the first time from 56Ni and its decays: 56Ni ! 56Co ! 56F∗ ! 56F (10) with lines at 847 keV and 1 288 keV (measured at γ-ray satellites). Neutrinos from SN1987A: a total of 4 detectors claimed to have seen neutrinos from 1987A: 2 water Cherenkov detectors (Kamiokande and Irwine-Michigan- Brookhaven, IMB) and 2 organic liquid scintillator detectors (Baksan and Mont Blanc). + • All detectors were sensitive to the inverse beta decay reaction νe + p ! n + e but only 3 of them measured the neutrinos in time coincidence (within their time uncertainty) 8 1 LECTURE 4: SUPERNOVA NEUTRINOS • Mont Blanc detected neutrinos 4.5 h before the others (5 events at the energy threshold of the detector). Therefore, it is likely that these neutrinos were not related to the SN explosion • Figure6 shows the 24 neutrinos measured in coincidence Figure 6.

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