Manual Stellar evolution and nucleosynthesis

Free download. Book file PDF easily for everyone and every device. You can download and read online Stellar evolution and nucleosynthesis file PDF Book only if you are registered here. And also you can download or read online all Book PDF file that related with Stellar evolution and nucleosynthesis book. Happy reading Stellar evolution and nucleosynthesis Bookeveryone. Download file Free Book PDF Stellar evolution and nucleosynthesis at Complete PDF Library. This Book have some digital formats such us :paperbook, ebook, kindle, epub, fb2 and another formats. Here is The CompletePDF Book Library. It's free to register here to get Book file PDF Stellar evolution and nucleosynthesis Pocket Guide.
Buy Stellar Evolution and Nucleosynthesis on fuwydusuni.tk ✓ FREE SHIPPING on qualified orders.
Table of contents

Ships with Tracking Number! May not contain Access Codes or Supplements. May be ex-library. Buy with confidence, excellent customer service!. Items related to Principles of Stellar Evolution and Nucleosynthesis. Principles of Stellar Evolution and Nucleosynthesis. Publisher: McGraw Hill Text , This specific ISBN edition is currently not available.

View all copies of this ISBN edition:. Synopsis About this title Donald D. From the Back Cover : Donald D. About the Author : Donald D. Clayton is professor of astrophysics at Clemson University. Buy Used Condition: Good Shows some signs of wear, and may Learn more about this copy. Other Popular Editions of the Same Title. Search for all books with this author and title. However, as shown in Ref. Therefore, such low Y e matter ejected due to multi-dimensional effects can explain the origin of weak-r-process elements.

Because of the lack of observations, it is not yet clear whether weak-r stars with enhancements of Sr, Y, Zr always have Nb—Mo enhancements as well. If this is the case, the innermost matter of hypernovae or the hot bubble of normal supernovae considered in Ref. It is currently not clear which astronomical site has such an environment. In this section we present remnant neutron star masses in both the UN and YU models and discuss their implications. Among them, the most important factor is the CO-core mass, which is determined by the stellar evolution before core collapse. A larger CO core leads to a larger remnant mass because it typically leads to a larger Fe core and, more importantly, it increases the amount of mass above the Fe core.

Once the pre-explosion density structure is given, the explosion energy determines the mass-cut or remnant neutron star mass M rem. For a given progenitor model, a larger explosion energy blows up more material above the Fe core, leading to a smaller mass-cut.

Stellar evolution and nucleosynthesis: the role of agb mass

In this paper, however, as well as in our previous work, we do not determine the mass-cut in this dynamical way but determine it by the amount of ejected 56 Ni, M 56 Ni. This is because M 56 Ni is rather sensitive to explosion energy when we determine the mass-cut dynamically. In reality, each SN may eject a somewhat different amount of 56 Ni. However, when we apply supernova yields, e. For this purpose, it is better to determine the mass-cut by M 56 Ni. Strictly speaking, to eject the same amount of 56 Ni, a more massive core requires a larger E exp.

However, the resultant M rem is not very different, so we fix the E exp for simplicity. It correlates with the CO-core mass, which in turn correlates with the He-core mass. M rem is the baryon mass and is not the observable neutron star mass. The conversion from baryon to gravitational mass depends on the EOS of nuclear matter.

Here we compare the M g of our models with observed neutron star NS masses. NS masses are most accurately determined when they are in double NS systems. In this case if two or more post-Keplerian parameters are obtained, NS masses are precisely determined e.


  • Navigation.
  • Voices of Alcoholism: The Healing Companion: Stories for Courage, Comfort and Strength (Voices Of series).
  • Haunted (Harlequin Presents);

Interestingly, the mass distribution shows a peak at 1. These stars may be considered to keep their birth mass [ 70 ], so they are suitable to compare with our results. NSs in binaries with high-mass companions are also considered to roughly keep their birth mass.

Reference [ 70 ] discussed the idea that the most likely values of the central mass and dispersions for these NSs are 1. NS masses in double NS systems and errors. As described above, this model may roughly correspond to the lightest Fe-core-collapse SNe. In the calculation in Ref. This seems to be larger than the observed minimum mass NS, though these O—Ne supernovae have not been well studied and currently only one progenitor model exists [ 71 , 72 ]; therefore, we do not know the general properties of these SNe yet. However, at this moment, we cannot say that the UN model represents reality better, because the number of double NS systems are still limited.

In Ref. Although this estimate is more uncertain than that in double NS systems, this wider range would be more easily understood in the YU model. Observations and theory are consistent in the sense that such massive NSs are rarely formed by normal supernova explosions. This explains why massive NSs are rarely formed. It will be very interesting to constrain observationally the method by which massive NSs can be formed at birth.

In this paper we have described our recent work on stellar evolution with some new results. Here we briefly review some other work by our group that is not mentioned above. Kuroda and Umeda have developed a 3D magneto-hydrodynamical general relativistic code with adaptive mesh refinements [ 75 ]. This code was used to follow gravitational Fe core collapse and to calculate the spectra of gravitational waves.

We will apply this code to various progenitor models to explore the explosion mechanism and nucleosynthesis. Okita and Umeda have developed a 2D special relativistic hydrodynamical code. Using this code, Okita and Umeda manuscript submitted for publication explored the conditions for successful ejection of ultra-relativistic jets in the collapsar model of GRBs. This code is also applied to explore various aspects of explosions and nucleosynthesis in SNe and GRBs. One example is given in the next subsection. SN bi was a very bright SN Ic and 3.

First, Gal-Yam et al. Unfortunately, without the light-curve data well before the maximum light, one cannot distinguish these two models. We should remember, though, that for the CCSN model we are assuming that such a large CO star can explode energetically. So far, no one has shown such a explosion from first-principles calculations. For spherically symmetric calculations, the same method as described in Sect.

For a jet-like explosion, the code mentioned in Sect. After the explosion simulations, nucleosynthesis was calculated post-processingly. First, we investigated the explosion-energy dependence of the 56 Ni amount. The ejected amount of 56 Ni is smaller than 2. The progenitor is a The formation of first-generation or Population Pop III stars in the Universe is considered to be quite different from that of later-generation stars. This is because, when the first-generation stars were formed, there was no metal, thus radiative feedback from the proto-star was weak and might not have been able to stop gas accretion.

Also, the gas accretion rate itself was larger than that in the present universe. In Ohkubo et al. Next, by taking account of a realistic mass accretion rate, Ref. These works, however, neglected the effects on the accretion disk likely formed around the first stars. First stars are formed around the center of a dark halo where the density of dark matter is much higher than in other places.

These stars typically have much larger radii, thus surface temperatures are lower and radiative feedback is expected to be weak. Umeda et al. It has been known that such huge stars are vibrationally unstable against the epsilon mechanism e. As long as the star is sustained by dark-matter annihilation and not nuclear burning, the epsilon mechanism does not operate.

However, after the main sequence stage, nuclear burning soon dominates unless the captured dark matter effect is important. Sonoi and Umeda [ 92 ] explored the stability of such very massive stars against the epsilon mechanism. Nozawa, T. Kozasa and collaborators. For example, Nozawa et al. The theory of dust formation was also applied to actual supernovae. It was also applied to SNe Ia [ 97 ]. These dust grains in supernovae play important roles in star and galaxy formation, and chemical evolution in galaxies.

Presolar grains are recovered from primitive meteorites or interplanetary dust and are identified as grains with very large isotopic anomalies compared with solar-system materials see, e. The observed isotopic ratios of the grains are considered to indicate traces of nucleosynthesis in stars at their birth or Galactic chemical evolution. Small amounts of presolar grains are considered to originate from supernovae. They mainly indicate excesses of 12 C, 15 N, and 28 Si. Some supernova grains show evidence for the original presence of radioactive 44 Ti in Ca isotopic ratios, which strongly supports their origin.

Isotopic ratios of heavy elements such as Mo and Ba have also been observed. However, it is still difficult to reproduce the observed isotopic ratios by supernova models. The bulk composition of supernova ejecta of supernova models does not reproduce the observed isotopic ratios of supernova grains. In order to reproduce the observed isotopic ratios, inhomogeneous mixing is required. Yoshida and Hashimoto [ ] and Yoshida [ ] investigated supernova mixtures, reproducing several isotopic ratios of supernova-originating SiC and graphite grains.

They divided supernova ejecta into seven different layers and investigated the mixing ratios, reproducing C, N, O, Al, Si, and Ti isotopic ratios as well as possible for individual grains. The mixing ratios of the mixtures strongly depend on the reproduced isotopic ratios. In this paper we have explained the differences between the newly developed efficient YU code and the previous UN code, and have shown that the YU code yields reasonable results, as shown in Sect. We need such an efficient code because the study of massive star evolution still requires a heavy amount of calculations, as described in the following.

Traditionally, the nucleosynthetic argument shown in Sect. However, as shown in Fig. As shown in Sect. One of our purposes in developing the efficient code was to tackle the evolutions just above and below Fe-core-forming SNe. As mentioned in Sect. It is not clear yet if such events cause mass loss and shock interactions to be observable.

In order to follow these stages precisely, it is important for a code to include the acceleration term. The stars just below the critical mass for Fe-core formation are very interesting for another reason.

Stellar Processes and Evolution

Such a star may become an electron capture supernova ECSN. Since these calculations require high numerical resolution in both time and space [ 8 ], we need an efficient code to perform the computation. We have also been working to develop a new code including the rotation effects as in Refs. This code is based on the YU code, and also aims to achieve efficient computation. Although rotating stellar models in the 1D formalism have been already calculated by these authors, angular momentum transfer is still quite uncertain.

Therefore, our understanding of the evolution of a rotating star is still far from complete. Since rotation is inevitable for constructing realistic progenitor models for hypernovae and GRBs, and possibly even for normal CCSNe, we plan to calculate rotating progenitor models as well. The uncertainties in the mass-loss rate are another reason why one needs to compute several cases to find a better set.

Since knowledge on the rate is still very limited, one needs to constrain the rate by using various pieces of information, including the properties of massive stars, supernovae, compact remnants, ISM abundances, and so on. Unfortunately, this may not be easy, because other uncertain factors may be involved, such as rotation, binarity, and magnetic field effects.

Supernovae may provide a key to understanding the rate. For example, as discussed in Ref. As for our nucleosynthetic work, in Sect. This model reasonably reproduces supernova yields up to Zn. This suggests that nucleosynthesis up to Zn does not greatly depend on the details of the central engine. On the other hand, nucleosynthesis of r- and weak-r-process elements depends on the details of the explosion model. As described in Ref.

These elements can be produced and ejected from the hot bubble of normal SNe. However, to produce heavier weak-r-process elements, Mo and Ru, a much larger entropy and a neutron-rich environment is required [ 56 ]. It is currently not certain if such an environment can be realized in a CCSN.

This is because long-term simulations of a CCSN have shown that high-entropy matter ejected from a CCSN is always almost neutral or proton-rich [ 61 ]. With this process, such weak-r-process elements may be synthesized in proton-rich matter if the entropy is sufficiently high.

Stellar Evolution and Nucleosynthesis | NHBS Academic & Professional Books

Thus it is critically important to observationally clarify if normal SNe produce Mo—Ru or not. We should note that the results in Ref. The main reason that the matter becomes proton-rich is because of the interaction between matter and neutrinos. If the explosion is driven with the assistance of something else, such as rotation energy, neutron-rich matter may be ejected. The degeneracy of electrons also increases in the contracting core and the above-mentioned temperature inversion appears. This ignition causes the core to expand adiabatically.

Account Options

After the burning terminates, the core contracts adiabatically again. After the second off-center Si ignition, the burning front reaches the center. The core transforms into an Fe core, then collapses. Because it has a larger C—O core of 1. Therefore, Ne and O ignite at the center. This is just because Ne has burned ahead of O, producing an O—Si core. Then off-center Si burning occurs Fig. In Fig. These stars lose all their H and He layers and become Wolf—Rayet stars. Metal-poorer stars have a similar M dependence of the final mass.

In this figure, the zero-age main-sequence mass, where the mass loss becomes effective, and the final mass for a given main-sequence mass become larger for metal-poorer stars. The final mass panel a , the He core mass panel b , and the CO core mass panel c with the relation to the main-sequence mass.

Top Authors

The surface He mass fraction of the stars is about 0. The metallicity dependence of the He core is small [Fig. The mass of the He-rich shell is 1. Since the fraction of the He layer is small, mass loss brings about the removal of the He layer rather than a reduction in the He core mass. On the other hand, the CO core mass can be more massive in metal-poorer stars. Nucleosynthesis in these massive stars occurs mainly in two stages. One is before core collapse and the other is during supernova explosion. It is well known that, inside these massive stars, nuclear fusion takes place up to Fe synthesis.

Before gravitational core collapse, Fe core is produced at the center and lighter elements form onion-like structures from inside to the surface Fig. The corresponding final masses are The mass fraction distribution in the H-rich envelope outside the range of each figure is the same as that at the outer edge of the figure.

Nucleosynthesis during a supernova explosion is usually called explosive nucleosynthesis. During the explosion, a shock wave propagates out of the Fe core to the stellar surface. Behind the shock wave, matter heats up and nuclear burning takes place. Explosive nucleosynthesis is important for the synthesis of Si and heavier elements e. In our previous work e. We inject thermal or kinetic energy below the mass cut or just above the Fe core to initiate supernova shock.

Stellar evolution

We use a 1D PPM code for the hydrodynamical calculations and solve small alpha-networks together to calculate nuclear energy generation. Then we calculate the detailed nucleosynthesis by post-processing by solving a large nuclear reaction network. Explosive nucleosynthesis should depend on how the star explodes, but we do not yet know how gravitational collapse leads to the explosion.

Therefore, there are still several proposals for successful supernova explosions. Fortunately, the properties of supernova shock outside the Fe core are almost independent of the explosion mechanism, because in most models the supernova shock gains energy inside the Fe core. Then, as far as spherical symmetry is assumed, explosion energy E is the only parameter that determines the properties of the supernova shock. For example, the kinetic to thermal energy ratio does not particularly affect the propagation of the shock outside the Fe core, because it quickly converges into the same solution.

For most cases, nucleosynthesis of Fe peak elements including Zn and lighter elements can be safely calculated with this method. For a given progenitor model and explosion energy E , the propagation of shock and thus the time evolution of density distribution is also determined uniquely. Since the shocked region cools roughly adiabatically at first, the radiation entropy is often convenient to specify the explosion. In such a region, Si mostly decomposes to alpha-particles at first; this is an endothermic reaction. Then, as the temperature decreases, the alpha-particles start to recombine and alpha-rich freezeout nucleosynthesis takes place.

This phase is exothermic. In this complete Si-burning region, 56 Ni is dominantly produced but Co, Ni, and Zn are also mostly produced here. This problem and its solution are described in Subsect. This region is called the incomplete Si-burning region. The main products here are Cr, Mn, and 56 Ni. In these regions, nuclear burning also produces some amount of energy. In summary, previous work has shown that explosive nucleosynthesis up to Zn can be described well by instant energy injection models, though hypernova models over-produce Fe in simple spherically symmetric models.

Of course, neutrino emission depends on the explosion model and thus we have to assume a model to include the effects see Ref. The neutrino spectra are assumed to obey a Fermi distribution with zero chemical potential. Some F is produced in the explosive He burning from 15 N. In the previous section, we mentioned that the carbon abundance after helium burning is important for the abundance of Ne to Ca. Here we look at the results of the UN and YU models more closely.

Red and blue lines indicate the yield ratios in the YU model and UN model, respectively. This is also seen in the figure. The UN yields calculated with the same parameter choices were applied to the Galactic chemical evolution model in Ref. It was shown that the model could reproduce the observed abundance pattern reasonably well. In the early universe, where the metal abundance in interstellar matter ISM was low, the Fe to H ratio was determined by the amount of Fe produced by a supernova over that of hydrogen swept by the supernova shock.

As mentioned above, however, there is one problem with this idea. Zn and Co are produced mainly in the complete Si-burning region while Fe or 56 Ni is produced in both the complete and incomplete Si-burning regions. Under the assumption of a spherical symmetric explosion, this means that the mass-cut has to be taken sufficiently deep, then the ejected mass of Fe or 56 Ni also increases. To quantify this problem, Ref. In the model, yields are calculated as follows.

First, the innermost matter is mixed between the mass coordinates M in and M out. For hypernova models, M out is set to the upper boundary of the incomplete Si-burning region, and M in is chosen sufficiently deep to eject Zn. This suggests that the explosion mechanism for supernovae and hypernovae is different. In such models, it is assumed that an unknown central engine ejects jets along the polar directions. These jets can blow up the entire star above the Si layer, but the explosive Si burning mainly takes place along the jet directions. As a result, the mass fraction of complete Si-burning products becomes smaller than the spherically symmetric models.

These conditions are difficult to satisfy for the hot bubble matter of supernovae, but relatively easy for the innermost ejecta of hypernovae. It is quite likely that, during explosion, supernovae produce elements heavier than Zn, because some EMP stars have r- and weak-r-process elements. This suggests that r-process nucleosynthesis is almost universal.

There are other classes of stars that show a greater abundance of weak r-process elements such as Sr, Y, Zr than the universal r-process pattern. Such low Y e matter cannot be ejected as long as we consider an exactly spherically symmetric explosion. However, as shown in Ref. Therefore, such low Y e matter ejected due to multi-dimensional effects can explain the origin of weak-r-process elements. Because of the lack of observations, it is not yet clear whether weak-r stars with enhancements of Sr, Y, Zr always have Nb—Mo enhancements as well.

If this is the case, the innermost matter of hypernovae or the hot bubble of normal supernovae considered in Ref. It is currently not clear which astronomical site has such an environment. In this section we present remnant neutron star masses in both the UN and YU models and discuss their implications. Among them, the most important factor is the CO-core mass, which is determined by the stellar evolution before core collapse. A larger CO core leads to a larger remnant mass because it typically leads to a larger Fe core and, more importantly, it increases the amount of mass above the Fe core.

Once the pre-explosion density structure is given, the explosion energy determines the mass-cut or remnant neutron star mass M rem. For a given progenitor model, a larger explosion energy blows up more material above the Fe core, leading to a smaller mass-cut. In this paper, however, as well as in our previous work, we do not determine the mass-cut in this dynamical way but determine it by the amount of ejected 56 Ni, M 56 Ni. This is because M 56 Ni is rather sensitive to explosion energy when we determine the mass-cut dynamically.

In reality, each SN may eject a somewhat different amount of 56 Ni. However, when we apply supernova yields, e. For this purpose, it is better to determine the mass-cut by M 56 Ni. Strictly speaking, to eject the same amount of 56 Ni, a more massive core requires a larger E exp. However, the resultant M rem is not very different, so we fix the E exp for simplicity. It correlates with the CO-core mass, which in turn correlates with the He-core mass. M rem is the baryon mass and is not the observable neutron star mass.

The conversion from baryon to gravitational mass depends on the EOS of nuclear matter. Here we compare the M g of our models with observed neutron star NS masses. NS masses are most accurately determined when they are in double NS systems. In this case if two or more post-Keplerian parameters are obtained, NS masses are precisely determined e.

Interestingly, the mass distribution shows a peak at 1. These stars may be considered to keep their birth mass [ 70 ], so they are suitable to compare with our results. NSs in binaries with high-mass companions are also considered to roughly keep their birth mass. Reference [ 70 ] discussed the idea that the most likely values of the central mass and dispersions for these NSs are 1. NS masses in double NS systems and errors.

As described above, this model may roughly correspond to the lightest Fe-core-collapse SNe. In the calculation in Ref.


  • Account Options;
  • Commemoration and Bloody Sunday: Pathways of Memory (Palgrave Macmillan Memory Studies)?
  • Reactions involving 12C: Nucleosynthesis and Stellar Evolution?
  • Top Authors.
  • Stellar Evolution and Nucleosynthesis!
  • How To Pay The Rent With Your Camera - THIS MONTH! (N/A)!
  • Principles of Stellar Evolution and Nucleosynthesis - Semantic Scholar?

This seems to be larger than the observed minimum mass NS, though these O—Ne supernovae have not been well studied and currently only one progenitor model exists [ 71 , 72 ]; therefore, we do not know the general properties of these SNe yet. However, at this moment, we cannot say that the UN model represents reality better, because the number of double NS systems are still limited.

In Ref. Although this estimate is more uncertain than that in double NS systems, this wider range would be more easily understood in the YU model. Observations and theory are consistent in the sense that such massive NSs are rarely formed by normal supernova explosions. This explains why massive NSs are rarely formed. It will be very interesting to constrain observationally the method by which massive NSs can be formed at birth.

In this paper we have described our recent work on stellar evolution with some new results. Here we briefly review some other work by our group that is not mentioned above. Kuroda and Umeda have developed a 3D magneto-hydrodynamical general relativistic code with adaptive mesh refinements [ 75 ]. This code was used to follow gravitational Fe core collapse and to calculate the spectra of gravitational waves.

We will apply this code to various progenitor models to explore the explosion mechanism and nucleosynthesis. Okita and Umeda have developed a 2D special relativistic hydrodynamical code. Using this code, Okita and Umeda manuscript submitted for publication explored the conditions for successful ejection of ultra-relativistic jets in the collapsar model of GRBs.

This code is also applied to explore various aspects of explosions and nucleosynthesis in SNe and GRBs. One example is given in the next subsection. SN bi was a very bright SN Ic and 3. First, Gal-Yam et al.