Physics 2018: Great Ideas in Science: The Physics Module Nuclear Physics Lecture Notes

Physics 2018: Great Ideas in Science: The Physics Module Nuclear Physics Lecture Notes Dr. Donald G. Luttermoser East Tennessee State University Edition 1.0

Abstract These class notes are designed for use of the instructor and students of the course Physics 2018: Great Ideas in Science. This edition was last modified for the Fall 2007 semester.

III. Nuclear Physics A. The Structure of Matter. 1. Matter exists in 3 different thermodynamic states: a) A solid is a rigid body takes a lot of energy to change its shape. Solids can be classified into two types: i) Crystalline solids have atoms that are structured in an orderly fashion. ii) Amorphous solids have randomly arranged atoms. b) A liquid is fluid in nature moderate energy required to change its shape. c) A gas is also fluid in nature little energy required to change its shape. 2. If a gas gets hot enough, electrons circling the nucleus of the atoms in the gas are ripped away from the nucleus = the gas becomes ionized = ionized gas is called a plasma. 3. Matter consists of a distribution of particles (atoms and molecules). a) Atoms consist of a nucleus surrounded by electrons (which are negatively [ ] charged). The nucleus consists of protons (positive [+] charge) and neutrons (no [0] charge). b) H, He, Li, Be, B, C, N, O, etc. are the elements of the periodic table = atoms. The number of protons in the nucleus defines each atom. c) For elements heavier than H, typically the number of neutrons is equal to the number of protons. Isotopes of III 1

atoms contain different numbers of neutrons in the nucleus (e.g., 12 C, 13 C, and 14 C are isotopes of carbon). d) Neutral Atoms: # protons = # electrons. i) If electrons are taken away from the atom such that the number of protons exceeds the number of electrons, the atom becomes a positive ion (e.g., H + = H II = singly ionized hydrogen). ii) If the number of electrons exceeds the number of protons in the nucleus, the atom becomes a negative ion (e.g., H = negative hydrogen ion, a hydrogen atom with two electrons instead of one). iii) The ionization stage of an atom can be labeled in a variety of different ways: = Roman numerals: I = neutral atom (e.g., He I, neutral helium), II = singly ionized (e.g., Fe II), III = doubly ionized (e.g., C III), etc. = + exponents (positive ions): if no exponent appears, then we have a neutral atom (e.g., He, neutral helium), + = singly ionized (e.g., Fe + ), 3+ = triply ionized (e.g., C 3+ ), etc. = exponents (negative ions): = one extra electron (e.g.,h ), = two extra electrons (e.g., C ), etc. e) Molecules are a collection of atoms that are bound together by molecular bonds. i) Salt: NaCl (1 sodium atom + 1 chlorine atom). III 2

ii) Water: H 2 O (2 hydrogen atoms + 1 oxygen atom). iii) Methane: CH 4 (1 carbon atom + 4 hydrogen atoms). f) Molecules can adhere to each other through chemical bonds making a structured lattice = solids. B. An Overview of Nuclear Reactions. 1. In order to get nuclear particles to interact with each other, one needs to collide these particles with a very high kinetic energy (i.e., high speed). a) Since the nucleus is composed of both neutrons (zero charge) and protons (positive charge), two nuclei close together will experience a strong repulsive force due to Coulomb s electric force law. b) Remember that forces are derivatives of potential energy. Hence if one wishes to bring two nuclei closer and closer together, one must increase the kinetic energy of said nuclei to high enough values to overcome the electric potential field. c) From thermodynamics heat is nothing more than the average kinetic energy of the particles that make up a system (i.e., matter). d) Hence, one way to get nuclei moving fast is to heat them up. 2. Reactions in chemistry and physics can be one of two types: a) Endothermic reactions absorb energy. III 3

b) Exothermic reactions release energy. 3. The element iron (Fe) has a nucleus that is the most stable of all the atomic nuclei. a) Elements lighter (i.e., less massive) than iron will produce exothermic nuclear reactions when they fuse together to make heavier elements = nuclear fusion. b) Elements heavier (i.e., more massive) than iron will produce exothermic nuclear reactions when they break apart to make lighter elements = nuclear fission. 4. The amount of energy released duing these exothermic reactions is proportional to the amount of the mass difference between the parent (original) and daughter (offspring) particles through Einstein s famous equation E = mc 2. (III-1) 5. Conservation Laws. a) If we assign a baryon number B of +1 to each baryon (nucleon or hyperon) and 1 to an antinucleon or antihyperon, then in a closed system B = constant. (III-2) b) Similarly, if we assign a lepton number L of +1 to each lepton (i.e., e,µ,ν, etc.) and of 1 to antileptons (i.e., e +, µ, ν, etc.), then in a closed system L = constant. (III-3) c) A similar conservation law does not exist for bosons the mesons and field particles (i.e., photons). III 4

d) Charge must be conserved in a nuclear reaction. e) Mass-energy, via E = mc 2, must be conserved in a nuclear reaction. f) Momentum must be conserved. Hence a matter-antimatter reaction must create two photons to conserve momentum (e.g., e + e + 2γ). C. Thermonuclear Reactions. 1. In 1938, it became clear that the long-term energy source for stars must be thermonuclear fusion reactions. In these reactions, lighter elements burn to form heavier elements = nucleosynthesis. 2. Two nuclei will fuse to form one nuclei if they come within 10 13 cm of each other but they must be moving fast enough to overcome the Coulomb repulsion that exists between like charged particles. a) Particles must be at a high temperature to be moving fast. b) This high temp completely ionizes all of the nuclei. c) Temps must build even more to get the kinetic energy to overcome the Coulomb barrier. 3. In main sequence stars, H is fused into He. Since H is composed of 1 baryon and He, 4 baryons (2p + 2n), 4 H nuclei must be used to construct one He nuclei. 4 m H = 4 1.0078 amu = 4.0312 amu m He = 4.0026 amu m = 0.0286 amu III 5

a) This mass deficit, m, is converted into energy: E = mc 2 = (0.0286) (1.66 10 24 gm) (9.00 10 20 cm 2 /s 2 ) = 4.3 10 5 erg b) From this calculation, we see that the energy release efficiency, η, of this H He reaction is 0.0286/4.0312 = 0.0071 = only 0.71% of the original mass of H is converted to energy! c) With this in mind, we can rewrite Einstein s famous equation as E = ηmc 2, (III-4) where η is the efficiency of the reaction and m is the initial mass in the reaction. d) We can calculate the total energy the Sun will release during its main sequence lifetime. Since the reactions are only being carried out in the core of the Sun and this core contains about 10% of the Sun s mass, the total energy release will be E tot = 0.1M ηc 2 = 1.28 10 51 ergs. e) The present luminosity of the Sun is 3.90 10 33 erg/s. If the Sun s luminosity remains somewhat constant while on the main sequence, we can determine its main sequence lifetime: t MS ( ) =E tot /L =3.28 10 17 sec = 1.04 10 10 yrs, the Sun s main sequence lifetime is about 10 billion years. Since the Sun is currently 5 billion years old, it is at midlife. III 6

4. As previously mentioned, high temperatures are needed to overcome the Coulomb repulsion of the charged nuclei. But how high a temperature? a) From classical physics, we can set the average kinetic energy of the particles involved equal to the thermal energy of the particles and solve for the temperature (see page 334 of Carroll and Ostlie). This gives T classical = 2Z 1Z 2 e 2, (III-5) 3k B r where Z 1 e is the charge on particle 1, Z 2 e is the charge on particle 2, and r is the distance where a nuclear reaction will occur (about the size of a nucleus, 10 13 cm). b) For 2 protons coming together (Z 1 = Z 2 = 1), this gives a temperature of 10 10 K, whereas the Sun s central temperature is only 1.58 10 7 K. c) We could also investigate this in terms of energy. For the Sun s central temperature, each proton will have a thermal energy of 1 kev, whereas the Coulomb potential barrier is 1000 kev (1 MeV)! Not all the particles have this energy, some are moving at much quicker velocities, hence have higher thermal energies and temperatures following the Maxwell-Boltzmann distribution of velocities. Unfortunately, the number of particles in the tail of this Maxwellian is insufficient to account for the Sun s luminosity. III 7

d) We can now turn to quantum mechanics to solve the problem. As discussed in the subsection on particle spin, in reality elementary particles are not little billiard balls colliding with each other as a result of following trajectories. Instead, they follow probability distributions described by their wave functions. In quantum mechanics, there is a small probability that wave functions can penetrate energy barriers that are higher than the energy of the wave function. This effect is known as quantum tunneling. e) Using quantum mechanics, we can describe a temperature needed to produce a sufficient number of tunneling events to sustain a nuclear reaction (see page 335 of the textbook) as T quantum = 4µ mz1 2Z2 2 e4 3k B h 2, (III-6) where µ m is the reduced mass of the colliding particles and h is Planck s constant. f) In this equation, two protons can come together (i.e., fuse) at a temperature of 10 7 K, which is consistent with the central temperature deduced for the Sun. g) A more detailed calculation from statistical mechanics shows that the bulk of the energy is being liberated by reactions involving particles in the high energy tail of a Maxwellian distribution. i) Particles with energies at the Gamow Peak will be the ones that supply most of the energy through thermonuclear reactions. ii) The Gamow Peak corresponds to a local maximum in the two probability functions: the e E/k BT III 8

Maxwell-Boltzmann distribution term and the e be 1/2 quantum tunneling penetration term, where, b 23/2 π 2 µ 1/2 m Z 1Z 2 e 2 h. iii) The Gamow Peak for a given temperature will occur at the energy of ( bkb T E = 2 )2/3 for the Sun, the Gamow Peak is at 6 kev., (III-7) 5. By making use of statistical mechanics in conjunction with quantum mechanics, stellar interior modelers set up power laws that describe the energy production rate per unit mass of the form ε = ε X i X x ρ α T β, (III-8) where the X s are the mass fractions of the fusing particles, and ε,α. and β are constants that depend upon the reactions involved (more to come on this). D. Various Reaction Chains 1. Two different fusion processes convert H into He, the first is important for stars with T c < 1.8 10 7 K(M < 1.3M, F5 V star) and is called the proton-proton chain. a) The first of this reaction chain is called the PP I chain: Energy Reaction Reaction Released Time (MeV) 1 H+ 1 H 2 H+e + + ν e 1.442 1.4 10 9 yr 1 H+ 2 H 3 He + γ 5.493 6 sec 3 He + 3 He 4 He + 1 H+ 1 H 12.859 10 6 yr III 9

i) 1 H = hydrogen atom (1 proton). ii) 2 H = heavy hydrogen (1 proton + 1 neutron) = deuterium. iii) 3 He = light helium (2 protons + 1 neutron). iv) 4 He = helium (2 protons + 2 neutron) = alpha particle. v) γ = Gamma ray photon. vi) e + = positron (positive charge) = anti-electron (antimatter). This positron interacts with the free electrons in the core virtually immediately which produces 2 additional gamma ray photons. vii) ν e = electron neutrino (neutral particle). The neutrino s absorption cross section is negligible and leaves the stellar core (and star) immediately without further interaction. The energy loss from the neutrino is 0.263 MeV which has not been included in the Energy Released column. viii) In this PP I change, please note that the following reaction can take place 1.4% of the time that the first reaction takes place 1 H+e + 1 H 2 H+ν e the so-called pep (proton-electron-proton) reaction which releases 1.4 MeV and loses an additional 1.4 MeV in energy loss from the neutrino. III 10

ix) Note that the Energy Released column is a combination of the energy of any gamma rays created, the kinetic energy = thermal energy picked up by the resulting nuclei as a result of this reaction, and the energy gained by the positron annihilation. x) The times listed for the Reaction Time column is that for the Sun s central temperature. xi) Note that the first two reaction steps must occur twice before the last one can take place. xii) The last step of this reaction chain is occurs 69% of the time in comparison to the other two PP chains in the production of 4 He in the Sun. xiii) This PP I chain dominates the other PP chains in stars with central temperatures of T < 1.6 10 7 K. xiv) Note the long average time it takes for the first and third reaction steps to take place for a single particle. However when normalized by the total number of particles in the Sun s (or star s) core, about 9.0 10 37 of these reactions take place per second! b) A second chain, called PP II, also can occur (31% of the time in the Sun) in the production of 4 He once the first two steps of the PP I chain occur. III 11

Energy Reaction Reaction Released Time (MeV) 3 He + 4 He 7 Be + γ 1.586 1.0 10 6 yr 7 Be + e 7 Li + ν e 0.861 0.4 yr 7 Li + 1 H 4 He + 4 He 17.347 10 4 yr i) 7 Be = beryllium-7 atom (4 protons + 3 neutrons). ii) 7 Li = lithium-7 atom (3 protons + 4 neutrons). iii) The neutrino energy loss in this PP chain is 0.80 MeV. iv) This PP II chain dominates the other PP chains in stars with central temperatures of 1.6 10 7 < T < 2.5 10 7 K. v) The Energy Released and the Reaction Time have the same meaning as they did for the PP I chain. c) A third and final PP chain, called PP III, occurs only 0.3% of the time in the Sun in the production of 4 He once the first two steps of the PP I chain occurs and the first step of the PP II chain occurs. Energy Reaction Reaction Released Time (MeV) 7 Be + 1 H 8 B+γ 0.135 70 yr 8 B 8 Be + e + + ν e 17.98 1 sec 8 Be 4 He + 4 He 0.095 1 sec III 12

i) 8 B = boron-8 atom (5 protons + 3 neutrons). ii) 8 Be = beryllium atom (4 protons + 4 neutrons). iii) The neutrino energy loss in this PP chain is 7.2 MeV. The Davis solar neutrino experiment, which detected only 1/3-rd of the predicted solar neutrinos was most sensitive to these 8 B beta decay neutrinos. Recently, neutrinos have been found to oscillate between the 3 known neutrino states which accounts for the low detection rate of the Davis experiment. iv) This PP III chain dominates the other PP chains in stars with central temperatures of T > 2.5 10 7 K in the production of helium. v) The Energy Released and the Reaction Time have the same meaning as they did for the PP I chain. 2. For more massive stars (T c > 1.8 10 7 K, M > 1.3M, F5 V star), the CNO cycle is the dominant reaction chain. a) This reaction chain uses carbon as a catalyst: Energy Reaction Reaction Released Time (MeV) 12 C+ 1 H 13 N+γ 1.95 1.3 10 7 yr 13 N 13 C+e + + ν e 2.22 7 min 13 C+ 1 H 14 N+γ 7.54 2.7 10 6 yr 14 N+ 1 H 15 O+γ 7.35 3.2 10 8 yr 15 O 15 N+e + + ν e 2.71 82 sec 15 N+ 1 H 12 C+ 4 He 4.96 1.1 10 5 yr i) 12 C = carbon-12 (6 protons + 6 neutrons). III 13

ii) 13 C = carbon-13 (6 protons + 7 neutrons). iii) 13 N = nitrogen-13 (7 protons + 6 neutrons) [radioactive]. iv) 14 N = nitrogen-14 (7 protons + 7 neutrons). v) 15 N = nitrogen-15 (7 protons + 8 neutrons). vi) 15 O = oxygen-15 (8 protons + 7 neutrons) [radioactive]. vii) 16 O = oxygen-16 (8 protons + 8 neutrons). viii) The neutrino energy loss in the 13 N beta decay is 0.710 MeV and the neutrino energy loss is 1.000 MeV for the second 15 O beta decay. ix) The Energy Released and the Reaction Time have the same meaning as they did for the PP I chain. b) Note that this reaction sequence does not make any new elements other than He! c) For the last step in the CNO cycle, an additional set of reactions can take place: III 14

Energy Reaction Reaction Released Time (MeV) 15 N+ 1 H 16 O+γ 12.126 1.0 10 7 yr 16 O+ 1 H 17 F+γ 0.601 3.0 10 10 yr 17 F 17 O+e + + ν 2.762 3 min 17 O+ 1 H 14 N+ 4 He 1.193 2.0 10 11 yr i) 17 F = fluorine-17 (9 protons + 8 neutrons). ii) 18 O = oxygen-16 (8 protons + 10 neutrons). iii) The neutrino energy loss in the 17 F beta decay is 0.94 MeV. iv) The resulting 14 N isotope can then be used back in the 14 N+ 1 H 15 O+γ reaction in the primary CNO cycle. 3. The thermonuclear reaction rate, ε (in erg/gm/s), is very sensitive to temperature. For the two hydrogen to helium reaction chains (i.e., the proton-proton chain and CNO cycle), we can write Eq. (IV-48) as two separate equations: ( T )β ε pp = ε ρx 2 10 6 (III-9) ( T )β ε cc = ε ρxx CN 10 6, (III-10) where X is the mass fraction of hydrogen (as defined on page IV-8 of the course notes), X CN is the weighted average of the combined mass fraction of carbon and nitrogen (since these species are the III 15

two lead-off species of the two CNO cycles), note that typically X CN = 1 3 Z, (III-11) with the Z being the metalicity mass fraction. Finally, ε and α are temperature dependent constants given in the following table (from Bosmas-Crespin, Fowler, Humblet 1954, Bull. Soc. Royale Sciences Liege, No. 9-10, 327). ε pp ε cc T/10 6 log ε α T/10 6 log ε α 4 6 6.84 6 12 16 22.2 20 6 10 6.04 5 16 24 19.8 18 9 13 5.56 4.5 21 31 17.1 16 11 17 5.02 4 24 36 15.6 15 16 24 4.40 3.5 36 50 12.5 13 In this table, for overlapping temperatures, a weighted average is used to get the final rate. 4. We will see shortly, that when a main sequence uses up all of its H fuel in its core, the now He-rich core will contract and heat up. When temperatures exceed 10 8 K, helium fusion can begin. Helium fuses via the triple-α process. The ash of this reaction is carbon. III 16

a) The reaction is as follows: Energy Reaction Released (MeV) 4 He + 4 He 8 Be + γ -0.0921 8 Be + 4 He 12 C+γ 7.37 i) Since this reaction chain is not occurring in the Sun, we have not reported on any reaction times here since they are very temperature dependent. ii) As can be seen, three α-particles (i.e., He nuclei) fuse to become one carbon nuclei. iii) 8 Be is unstable and quickly decays, so there is not much berylium around for the 2nd chain to take place. For every 1 berylium nuclei, there are 10 10 α-particles, however this ratio is more than enough to release enough energy to power a red giant star. Note that this reaction actually drains energy away from the gas = it is an endothermic reaction. All reactions we have mentioned up to now have been exothermic = they release energy. b) The 3α process has an even greater sensitivity to temperature than the CNO cycle: ( T )30 ε 3α 10 8 ρy 3 10 6 erg/gm/s, (III-12) where Y is the fractional abundance of helium per unit mass. c) This is the way the Universe makes carbon. As such, the C atoms that make up our DNA were created in an III 17

ancient red giant star that no longer exists. To quote Carl Sagan, we are star stuff! d) Should the core of a red giant obtain temperatures that exceed a few hundred million Kelvins, another reaction can take place via an alpha (α) capture: 12 C+ 4 He 16 O+γ which releases 7.161 MeV of energy. e) Most of the 16 O in the Universe is made in this fashion. 5. Finally, if somewhat higher temperatures are ever encountered inside a star, which happens during stellar evolution of massive stars, even heavier elements can be created from the fusion of additional α-particles and α-particle by-products: Energy Minimum Reaction Released Temperature (MeV) Required (10 6 K) 14 O+ 4 He 20 Ne + γ 4.730 700 20 Ne + 4 He 24 Mg + γ 9.317 1500 24 Mg + 4 He 28 Si + γ 9.981 1800 28 Si + 4 He 32 S+γ 6.948 2500 32 S+ 4 He 36 Ar + γ 6.645 3500 12 C+ 12 C 24 Mg + γ 13.930 800 16 O+ 16 O 32 S+γ 16.539 2000 a) Various silicon burning reactions can occur at temperatures exceeding 3 10 9 K. Silicon burning produces the iron (Fe) group elements. b) Once Fe is formed, reactions that produce heavier elements are all endothemic and have a tough time forming via the standard thermonuclear burning. Such elements, III 18

and the elements not built upon α-particles, are created via the r- (for rapid neutron capture) and s- (for slow neutron capture) processes. These processes will be discussed in the supernovae section of the course notes (i.e., VII). c) The reaction times of this heavy element nucleosynthesis will be discussed in the stellar evolution sections of the notes. III 19