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Synthesis of heavy elements

Nucleosynthesis paths
Fig. 1: Nucleosynthesis paths of the various processes.

The chemical elements heavier than iron can not be produced by fusion reactions. Since the binding energy per nucleon is decreasing past the iron group nuclei and because of the increasing Coulomb barriers, the nuclei heavier than iron are essentially synthesised by neutron capture reactions. Three different processes can be identified in this mass region, the slow (s) neutron capture process during stellar He burning, the rapid (r) neutron capture process, and the photo-dissociation (p) process, the latter two being presumably related to Supernova explosions. Whether a fourth process, the rapid proton (rp) capture process, contributes also to the production of the heavy elements or whether further additional processes are needed to explain their abundances is still under discussion. A comprehensive and quantitative description of these processes is mandatory for the understanding of stellar evolution, of stellar explosion mechanisms, and of galactic chemical evolution. Isotopic abundance patterns of the elements between Fe and the actinides found in the solar system or in meteorites in form of presolar grains carry important information about the stellar production mechanisms and even about the physical conditions at the respective stellar sites. The interpretation of the observed abundance patterns by sophisticated models of stellar evolution can be used as a crucial test for the proper description of the related nucleosynthesis processes.

The s-process

Planetary nebula
Fig. 2: Planetary nebula NGC 3132at. At the end of their life, Red Giant stars expel the outer envelope and form a planetary nebula. In this way, the freshly synthesized s-process material is mixed back into the interstellar gas. The star itself ends as a white dwarf which can be seen in the center of the nebula.
(Credit: The Hubble Heritage Team)

About half of the elemental abundances between Fe and Bi are produced by the s process, which is associated with stellar He burning scenarios of evolved Red Giant stars [1]. Since the s process is characterized by relatively low neutron densities neutron-capture times are in general much longer than typical β-decay half-lives. Consequently, the s process works its way along the valley of beta stability, starting at the abundant seed nuclei of the iron group elements and ending at the alpha-unstable trans-bismuth isotopes. The abundances of the so produced isotopes are inversely proportional to their neutron capture cross sections. The most important ingredients for a quantitative description of the s process are neutron capture cross sections and β-decay half-lives of the involved nuclei. However, nuclear properties, such as lifetimes, in the hot stellar plasmas can differ from the values measured in the laboratory. This was demonstrated by two experiments at GSI which investigated the bound-state beta decay of 163Dy and 187Re. In both cases it was experimentally proven that the lifetime in hot stellar plasmas, where the atoms are fully ionized, is considerably shorter than the terrestrial value (163Dy is even stable under terrestrial conditions!). This has consequences on the s-process path.


[1] Busso, M., Gallino, R., and Wasserburg, G.J., Ann. Rev. Astron. Astrophys. 37, 1999, 239.

The r-process

Kepler's supernova
Fig. 3: Remnant of Kepler's supernova.
(Credit: NASA, ESA)

The r process, which is responsible for the production of about half of the heavy-element abundances including Th and U, is characterized by enormous neutron densities of 1020 to 1024 cm-3 and time scales of a few seconds. These conditions clearly point to an explosive scenario, e.g. core collapse Supernovae (SNII), for the astrophysical site of the r process. Due to the high neutron densities, the seed nuclei capture many neutrons in a very short time interval, driving the r-process path along the isotopic chains to very neutron-rich species until (n,γ)/(γ,n) equilibrium is established at the so-called waiting points. After the β-decay to the next higher element the sequence of neutron captures continues until a new equilibrium is established. During the peak neutron flux, the relevant nuclear physics information concerns the neutron separation energies, which determine the (n,γ)/(γ,n) equilibrium, and the β-decay rates of the waiting point nuclei, which are important for the duration of the r-process and for the r-process abundance pattern (Nr~1/t½). Therefore, mainly masses and β-half lives of the very neutron-rich nuclei on the r-process path are required for obtaining the primary r-process yields.

Role of fission in r-process nucleosynthesis

Fission can have an important influence on the termination of the r-process and on the abundances of long-lived actinides, which are relevant for determining the age of the Universe. Fission can also influence the abundances of nuclei in the region A~90 and 130 due to fission cycling. In order to quantitatively understand the role of fission in the r process, two important pieces of information are needed: the fission-barrier heights and mass- and charge-distributions of the fission fragments.

Since experimental information is only available for nuclei in a limited region of the nuclide chart, the heavy r-process nuclei. have to be described by theoretical predictions. Recently, important progress has been made in developing full microscopic approaches to nuclear fission. Nevertheless, due to the complexity of the problem, this type of calculations is still difficult to apply to heavy nuclei and the accuracy of these models is rather limited. At GSI a macroscopic-microscopic approach was used to investigate the fission contribution to the r process. A model for calculating mass- and charge-distributions of fission fragments that can correctly predict the transition from double-humped to single-humped distributions with decreasing mass of the fissioning system and increasing excitation energy in the light actinides was developed and has been used [2] to calculate fission-fragment distributions in neutrino-induced fission of r-process nuclei.


[1] A. Kelic and K.-H. Schmidt, Phys. Lett. B 634 (2006) 362
[2] A. Kelic et al, Phys. Lett. B 616 (2005) 48

The p-process

There are 32 proton rich nuclei between Se and Hg, which can neither be produced by the s process nor by the r process. These nuclei are attributed to the p process, which requires high temperatures of about 2-3 GK. In such an environment the reaction flow is carried by photo-dissociation processes, i.e. by the (g,n), (g,p), and (g,α) channels. Since high temperatures are needed, the presently favoured sites for the p process are the explosively burning O/Ne layers in Supernovae of type II, where temperatures of 2 - 3 109 K are maintained for about 1 s at densities of ~106 g cm-3 [1,2]. Under these conditions, proton-rich nuclei are produced by a sequence of (g,n) reactions. When this sequence is halted after about five steps by the increasing neutron-separation energies, the further reaction flow is dominated by (g,p) and (g,a) reactions. As the temperature decreases after the explosion, the reaction path moves back to the region of stable nuclei. This scenario involves about 2000 nuclei connected by more than 20000 reactions and requires correspondingly large reaction networks to describe the abundance distributions following from these scenarios.


[1] Lambert, D.L., Astron. Astrophys. Rev. 3, 1992, 201
[2] M. Arnould and S. Goriely, Phys. Rep. 384, 2003, 1

The rp-process

The rapid proton capture (rp) process is the dominant nucleosynthesis and energy production process in explosive hydrogen burning scenarios. Starting from the break-out of the hot CNO cycle heavier nuclei are built by a sequence of proton captures and beta decays. This process plays an important role for X-ray bursters, which consist of a neutron star and a close companion. When hydrogen and helium rich material from the companion star is accreted onto the surface of the neutron star, explosive hydrogen burning is eventually triggered at critical temperatures and densities.

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