Nucleosynthesis Essay Example For Students

Nucleosynthesis Essay The big bang which created the universe, only created the elements Hydrogen (H) and Helium (He) and possibly a very small amount of Lithium (Li). However, a glance at the periodic table of the elements shows that today (some 15 billion years after the big bang) there are at least 108 known elements. Every atom of every element heavier than Li has been produced since the big bang! The factories which make these elements are stars. Nucleosynthesis or the synthesis of nuclei, is the process by which stars (which start out consisting mostly of H and He) produce all other elements. The key is nuclear fusion, in which small nuclei are joined together to form a larger nucleus. (This contrasts with nuclear fission, in which a large nucleus breaks apart to form two smaller nuclei). Fusion requires an extremely large amount of energy (see fig. 1), and can typically only take place in the centers of stars. FIGURE 1a) Low energy proton is strongly repelled by the 7Be nucleus.b) High energy proton moves so fast that it can strike the 7Be nucleus. Once the proton touches the nucleus, it has a chance to stick. If the proton sticks, the 7Be becomes a 8B nucleus.c) 8B is radioactive and changes into 8Be plus a positron (b+) and a neutrino (n). 8Be is itself radioactive, and almost immediately breaks into two 4He nuclei. Protons repel each other. This repulsion becomes stronger as the protons get closer together (just like when you try to stick two magnets together north to north, or south to south. Try this! As you push the magnets closer together, it becomes harder to do). However, if the protons can actually touch each other, they have a chance to stick together! This is because of the strong nuclear force which attracts nucleons (protons or neutrons) together, and is much stronger (at close range) than the electromagnetic force repulsion that makes protons repel other protons. (Magnets do not do this: two like poles will never stick together). In order to get a proton to strike another proton (or a nucleus that contains several protons) they must be traveling at high relative speeds; if their closing velocity is not great enough, they will never get close enough to stick together, because they strongly repel each other. But, just as you can make two of the same magnetic poles touch each other by providing sufficient force, so too can protons touch when they have sufficient relative speed. This can take place in the center of the sun, where the temperature is extremely high. Temperature is related to atomic motion: the hotter something is, the faster its atoms are moving see demo food coloring in water. Table 1 shows the nuclear reactions that are taking place in our sun, as well as nuclear reactions that take place in stars that are either older than our sun, or hotter than our sun. The reactions in columns 2 and 3 occur after a star has entered the red giant phase. How fast a star evolves to this point depends on its mass: stars heavier than the sun can reach this phase in less than 5 billion years (the age of the sun) whereas stars with about our suns mass take about 10 billion years to get there. The particles you may be unfamiliar with are: n the neutrino, g a gamma ray (high energy light wave), and b+ the positron (the antimatter version of the electron). TABLE 1. NUCLEAR REACTIONS IN STARSOUR SUN NOW OLDER, OR HOTTER STARSp + p 2H + b+ + n 4He + 4He 8Be + g 12C + p 13N + g2H + p 3He + g 8Be + 4He 12C + g 13N 13C + b+ + n3He + 3He 4He + p + p 12C + 4He 16O + g T1/2 = 10 min16O + 4He 20Ne + g 13C + p 14N + g3He + 4He 7Be + g 20Ne + 4He 24Mg + g 14N + p 15O + g7Be + p 8B + g 15O 15N + b+ + n8B 8Be + b+ + n T1/2 = 120 ms8Be 4He + 4He 15N + p 12C + 4HeHe burning (core) H burning shellIn our sun, the first three nuclear reactions shaded} are the major source of energy. The second group (of four) reactions also occur in the sun, but much less frequently than the first group (which is called the p-p chain). In both cases, the fuel (hydrogen) is converted into the product (helium), and energy (in the form of heat and light) is produced. The third column of reactions is called the CNO cycle, because carbon (C), nitrogen (N) and oxygen (O) are produced and C is recycled. The CNO reaction cycle is now occurring in the sun (the energy required fo r these reactions is roughly the same as that required for the p-p chain) because the sun had carbon to begin with (it hasnt made any C yet!). Since the sun had carbon present when it formed, it is referred to as a later generation star. The first generation of stars contained only H, He and Li from the big bang. Later generation stars contain material that has been processed in other stars. As a star like the sun evolves, a vast amount of H is consumed, and a vast amount of 4He is produced. This 4He can not combine with other 4He because this reaction requires more energy than is available, i.e. the 4He nuclei are not moving fast enough. However, as the 4He accumulates in the stars core, the pressure rises (which causes the temperature to rise since the core consists of gas). With increased temperature, more energy becomes available, and the star eventually reaches a point where the second column of above reactions can occur; this is called He burning. As we proceed down the second column, the reactions become increasingly less likely since they require increasing amounts of energy. This is because all nuclei have a positive electric charge, and like charges repel each other. The bigger the charge, the stronger the repulsion. This set of reactions can not produce any nucleus heavier than Fe (iron, atomic number 26). Again, a brief glance at the periodic table reveals th at there are many elements heavier than Fe; these are also produced in stars, but not by any kind of fusion reaction. What takes us beyond Fe are two nucleosynthetic processes, called the s-process and the r-process. Summary of â€Å"The Boston Photographs† Sample EssayThe word rapid is actually an understatement; it could be called explosive; the r-process occurs in supernova explosions! Heres how it works: Before a supernova, a star has produced an excessive amount of 56Fe. This accumulates in the core (recall that we cant go beyond 56Fe with fusion). As always, there is a battle between gravity (which tries to compact the core) and heat (which tries to expand the core). Eventually, after enough 56Fe is produced, gravity wins. When the Fe core collapses, it does so dramatically, and generates pressures which are truly incredible. The pressure is so great that the orbital electrons are pushed into their nucleus! Thus in one incredible electron capture reaction, all of the Fe in the core is converted to neutrons (1.4 solar masses worth!). An implosion shock wave reaches the cores center and rebounds. As it does so, it sweeps vast numbers of neutrons out with it, and they smash into the m atter above them. Now the neutron absorptions occur rapidly enough to bridge the gap out to nuclides like 70Zn. In fact, they happen rapidly enough to bridge the gap from Bi to the heavier elements (we know that 244Pu existed in the early solar system. This means that a 209Bi nucleus would have to absorb at least 35 neutrons before any a decay could occur!). Thus in one brief event, lasting at most only a few seconds, we produce all known elements heavier than Fe. Appendix (applications)There is an interesting application of s-process branching, through which we can roughly calculate the temperature that existed inside a red giant star even though the star exploded almost five billion years ago!The key is that dust is produced in the atmospheres of red giant stars, and the unique isotopic distributions of the elements made in the star are frozen in as solids. These dust grains survive to the present day, preserved in primitive meteorites see Interstellar Grains module. Lets look at the specific example of 85Kr to see how this remote thermometer works. Our bodies, at a temperature of about 40 C (100 F) give off infrared radiation which can be seen with special cameras. A log in a fire, at a temperature of about 600 C (1100 F) glows red. Molten metal in a furnace, at a temperature of about 1500 C (2700 F) shines with intense white light. Thus as temperature increases, the radiation (light) emitted becomes more energetic (changes color to shorter wavelengths) as well as more intense (more photons emitted per second). This is basically a result of the increased energy of the atomic collisions in the hot material see Blackbody Radiation module. For temperatures characteristic of star cores (hundreds of millions of C) the collisions produce nuclear reactions as well as an abundant supply of high energy gamma rays. When these gammas are absorbed by a nucleus, they can make the nucleus transition to an excited energy state (just as visible or ultraviolet light can make an atomic electron transition to a higher orbital. This is the first s tep in making a laser beam). As we saw in figure 1 of the Radioactive Decay module, 85Kr has a metastable excited state which is only 0.305 MeV above the ground state (a fairly small energy when considering nuclear transitions). The temperature in the star will dictate how much of the 85Kr present will be in its excited state (the temperature determines the number of photons and their energy distribution. This together with the amount of 85Kr nuclei present in the star (which can be roughly calculated) gives the amount of 85Kr nuclei which should be in the first excited state.) But from figure 2 of the Radioactive Decay module, we see that this excited state has a much shorter half-life than the ground state (4.48 hours vs. 10.7 years) and that this excited state can decay directly to 85Rb. Thus the more 85Kr that can reach this excited state, the shorter its effective half-life will be. Finally figure 2 of this module shows that 85Kr is a branch point on the s-process path. The 10.7 year half-life of 85Kr is su fficiently long that many nuclei will absorb a neutron to become 86Kr. However, a half-life of 4.48 hours is not long enough to absorb another neutron before beta decay (which happens 79% of the time from this excited state) and will not produce 86Kr. Therefore, the amount of 86Kr present in a dust grain tells us what percentage of the 85Kr absorbed a neutron, which in turn tells us the temperature that existed inside the star. This is a difficult idea. If you understood it, then you really understood this module