What is the compositon of a white dwarf or neutron star and why does it still shine?

Answer 1

White dwarf stars are made of degenerate gases. Stars less than half that of the sun become helium dwarfs (made up of helium-4 nuclei). Stars like our sun become carbon/oxygen dwarfs, since more massive stars can burn heavier fuels. Stars slightly more massive than the sun, but below 8 to 10 solar masses, can create neon/magnesium dwarfs. Although mainly gas, these compact bodies are quite dense (1 teaspoon would weigh several tons) due too extreme compaction of material.

Neutron stars (including pulsars) are formed from the remnants of super-massive stars. According to current theoretical models, the surface of a neutron star consists of an atmosphere about one meter thick, below which is a 'crust' of atomic nuclei and electrons a mile thick. Below this crust lies nuclei with increasing numbers of neutrons, until at a certain point below the surface (the neutron drip), free neutrons leak out of the nuclei. From this point to the core, is a 'slush' of neutrons, free electrons, and increasingly smaller atomic nuclei. Neutron-degenerate matter dominates the core, presumably in a super-fluid mixture with a few protons, neutrons, and other strange matter.

White dwarfs are essentially super hot cores...after an old star becomes a white dwarf, a white dwarf will radiate its stored heat out into space, like an ember. Over the course of several billion years, it will become a black dwarf...a dark ball of superdense gas in space.

Visible neutron stars (IE pulsars) are able to emit radiation through their rotational energy. Their rotational energy creates large magnetic fields, which excites particles into emitting photons. Neutron stars can also emit large bursts of radiation due to in-fall of matter from companion stars or the interstellar medium.

Answer 2

White Dwarfs Contain About The Same Amount Of Matter As The Sun.

Neutron Star Contain A Solid Crust With Iron And Similar Elemants

Answer 3

White dwarfs are the cores of stars that have died; The core of our sun will one day become a white dwarf, and it'll be very, very hot for a long time - it'll take a LONG time for the amount of heat & internal energy to finally bleed off into space, mostly because there's a LOT of heat, and very little surface area to radiate it.

Neutron stars are the cores of stars that have gone through a supernova. The electrons and protons have been crushed into each other, leaving a mass of nothing but neutrons. Neutrons have no charge, so they can sit right next to each other - eliminating ALL space in between the particles. (Your average atom is 99.9999999% empty space.) A teaspoon of neutron star would have the same mass as Mt. Everest.

Answer 4

Hi. A white dwarf and a neutron star are too separate things. The dwarf is still normal matter while the neutron star has been crushed by it's own gravity until the electrons touch the protons and merge into neutrons. Probably the hardest substance in the universe, my opinion. The dwarf is still a gas.

Answer 5

A neutron star is so named because it contains mostly neutrons, it still shines because although it no longer has fusion at it's core the amount of residual heat is so great it will radiate this energy for millions of years.

Answer 6

A neutron star is just that, neutrons. It shines because it's still hot from previous fusion.

Answer 7

Probably all stars with initial masses up to about eight solar masses finally end up as white dwarfs. Stars with more than about 8 solar masses explode as type II supernovae after a lifetime of only a few million years and become neutron stars or black holes. 90% of all stars finally become white dwarfs when their nuclear energy generation has ceased. With typically 0.6 solar masses and radii of about 10^9 cm (0.01 solar radii) the mean densities of white dwarfs are of the order of 10^5-10^6 g/cm^3 so that these stars can be considered as laboratories for matter at extreme densities and pressures. White dwarfs are stabilized against a gravitational collapse by the pressure of the degenerate electron gas in the interior. If the central density rho rises because of an increasing mass of the star, the electrons become relativistic, since the Pauli exclusion principle forces the majority of the electrons to reach relativistic velocities, even at low temperatures; by this effect the equation of state becomes "soft" (P~rho^(4/3)). The consequence is that no white dwarfs above the Chandrasekhar limiting mass of about 1.4 solar masses can be stabilized against a gravitational collapse.

Investigation of white dwarfs provide important constraints on the theory of stellar evolution. E.g. the existence of white dwarfs in young stellar clusters (with turn-off masses of about 5 solar masses ) show that massive stars must loose the bulk of their matter during their lifetime until they become a white dwarf with about 1.15 solar masses. The majority of white dwarfs have masses of about 0.6 solar masses and stems from progenitors with about 2 solar masses. The mass-loss processes during the stellar evolution is not well understood; the largest quantity is lost on the asymptotic giant branch (AGB) when the energy in a star is produced by alternating ignitions of a hydrogen and helium burning shell ("thermal pulses") in the interior. At the end of this phase about 0.1-0.3 solar mases are lost when the outer layers are blown away; this matter subsequently becomes visible as a planetary nebula. The nebula emits light since it is ionized by the hot (typically 50-200 kK) remnant. This central star of a planetary nebula is the direct progenitor of a white dwarf. After about 30000 years the nebula becomes invisible and the white dwarf cools down before they eventually become invisible. The coolest white dwarfs known have an age of about 9-10 billion years; the existence of these stars is a measure for the age of the disc of our galaxy.

With the exception of very cool objects the equation of state in the atmospheres of white dwarfs is the one for an ideal gas. For a white dwarf with an effective temperature of 25000K the particle densities and the temperatures range from 10^15-10^17 cm^-3 and 15-50 kK, respectively.

The chemical composition of white dwarfs is rather peculiar compared to main sequence stars: Normally only one element shows up in the optical spectra. The reason for this almost mono-elemental composition has been known for a long time. In 1949 Schatzman has shown that the huge gravitational acceleration (typically 10^8 cm sec^-2) together with the electric field leads to a downward diffusion of the heavier elements on time scales which are rather short compared to the evolutionary time scale.

The stars with spectral type DA are defined by the fact that their optical spectra exhibit the Balmer lines of hydrogen only. About 80% of all known white dwarfs belong to this class. Nearly all atmospheres of the non-DAs are dominated by helium. Depending on the effective temperature He II (spectral type DO, 45-120000K) or HeI (type DB, 11-28000K) lines are the strongest. On the cool end of the cooling sequence (below 11000K) the high energy needed to excite spectral lines leads to a continuous spectrum in the optical (type DC). In some stars with helium atmospheres, lines of carbon (type DQ) or other metals (Ca, Mg, Fe: spectral type DZ) show up in the spectra, probably due to accretion from the interstellar medium.

While many details and physical processes in the atmospheres of white dwarfs are well understood, the most obvious observational fact, namely that the white dwarfs are basically divided into two major groups --- the DAs and the non-DAs ---, is still unexplained. One possible explanation is that the spectral type is determined by the pre-white dwarf evolution. It has been suggested that the exact phase, when the star leaves the AGB, or a late thermal pulse in a post-AGB star determines whether significant amounts of hydrogen are left over in the atmosphere or not.

It is, however, very certain that transitions between the DA and the non-DA sequence must occur: between 28000K and 45000K not a single non-DA star exists. Since there is no physical process known by which the effective temperature of a DO with 45000K should jump down to 28000K the DO stars must become hydrogen rich at about 45kK for some, yet unknown reason.

First, the typical mass of a neutron star is about 1.4 solar masses, and the radius is probably about 10 km. By the way, the "mass" here is the gravitational mass (i.e., what you'd put into Kepler's laws for a satellite orbiting far away). This is distinct from the baryonic mass, which is what you'd get if you took every particle from a neutron star and weighed it on a distant scale. Because the gravitational redshift of a neutron star is so great, the gravitational mass is about 20% lower than the baryonic mass.

Anyway, imagine starting at the surface of a neutron star and burrowing your way down. The surface gravity is about 10^11 times Earth's, and the magnetic field is about 10^12 Gauss, which is enough to completely mess up atomic structure: for example, the ground state binding energy of hydrogen rises to 160 eV in a 10^12 Gauss field, versus 13.6 eV in no field. In the atmosphere and upper crust, you have lots of nuclei, so it isn't primarily neutrons yet. At the top of the crust, the nuclei are mostly iron 56 and lighter elements, but deeper down the pressure is high enough that the equilibrium atomic weights rise, so you might find Z=40, A=120 elements eventually. At densities of 10^6 g/cm^3 the electrons become degenerate, meaning that electrical and thermal conductivities are huge because the electrons can travel great distances before interacting.

Deeper yet, at a density around 4x10^11 g/cm^3, you reach the "neutron drip" layer. At this layer, it becomes energetically favorable for neutrons to float out of the nuclei and move freely around, so the neutrons "drip" out. Even further down, you mainly have free neutrons, with a 5%-10% sprinkling of protons and electrons. As the density increases, you find what has been dubbed the "pasta-antipasta" sequence. At relatively low (about 10^12 g/cm^3) densities, the nucleons are spread out like meatballs that are relatively far from each other. At higher densities, the nucleons merge to form spaghetti-like strands, and at even higher densities the nucleons look like sheets (such as lasagna). Increasing the density further brings a reversal of the above sequence, where you mainly have nucleons but the holes form (in order of increasing density) anti-lasagna, anti-spaghetti, and anti-meatballs (also called Swiss cheese).

When the density exceeds the nuclear density 2.8x10^14 g/cm^3 by a factor of 2 or 3, really exotic stuff might be able to form, like pion condensates, lambda hyperons, delta isobars, and quark-gluon plasmas.

At the moment of a neutron star's birth, the nucleons that compose it have energies characteristic of free fall, which is to say about 100 MeV per nucleon. That translates to 10^12 K or so. The star cools off very quickly, though, by neutrino emission, so that within a couple of seconds the temperature is below 10^11 K and falling fast. In this early stage of a neutron star's life neutrinos are produced copiously, and since if the neutrinos have energies less than about 10 MeV they sail right through the neutron star without interacting, they act as a wonderful heat sink. Early on, the easiest way to produce neutrinos is via the so-called "URCA" processes: n->p+e+(nu) [where (nu) means an antineutrino] and p+e->n+nu. If the core is composed of only "ordinary" matter (neutrons, protons, and electrons), then when the temperature drops below about 10^9 K all particles are degenerate and there are so many more neutrons than protons or electrons that the URCA processes don't conserve momentum, so a bystander particle is required, leading to the "modified URCA" processes n+n->n+p+e+(nu) and n+p+e->n+n+nu. The power lost from the neutron stars to neutrinos due to the modified URCA processes goes like T^8, so as the star cools down the emission in neutrinos drops sharply.

When the temperature has dropped far enough (probably between 10 and 10,000 years after the birth of the neutron star), processes less sensitive to the temperature take over. One example is standard thermal photon cooling, which has a power proportional to T^4. Another example is thermal pair bremsstrahlung in the crust, where an electron passes by a nucleus and, instead of emitting a single photon as in standard bremsstrahlung, emits a neutrino-antineutrino pair. This has a power that goes like T^6, but its importance is uncertain. In any case, the qualitative picture of "standard cooling" that has emerged is that the star first cools by URCA processes, then by modified URCA, then by neutrino pair bremsstrahlung, then by thermal photon emission. In such a picture, a 1,000 year old neutron star (like the Crab pulsar) would have a surface temperature of a few million degrees Kelvin.

But it may not be that simple...

Near the center of a neutron star, depending on the equation of state the density can get up to several times nuclear density. This is a regime that we can't explore on Earth, because the core temperatures of 10^9 K that are probably typical of young neutron stars are actually cold by nuclear standards, since in accelerators when high densities are produced it's always by smashing together particles with high Lorentz factors. Here, the thermal energies of the particles are much less than their rest masses. Anyway, that leaves us with only theoretical predictions, which (as you might expect given the lack of data to guide us) vary a lot. Some people think that strange matter, pion condensates, lambda hyperons, delta isobars, or free quark matter might form under those conditions, and it seems to be a general rule that no matter what the weird stuff is, if you have exotic matter then neutrino cooling processes proportional to T^6 can exist, which would mean that the star would cool off much faster than you thought. It even appears possible in some equations of state that the proton and electron fraction in the core may be high enough that the URCA process can operate, which would really cool things down in a hurry. Adding to the complication is that the neutrons probably form a superfluid (along with the protons forming a superconductor!), and depending on the critical temperature some of the cooling processes may get cut off.

So how do we test all this? We expect that after a hundred years or so the core will become isothermal (because it is then superfluid), and we can estimate thermal conductivities in the crust, so if we could measure the surface temperatures of many neutron stars, then we could estimate their core temperatures, which combined with age estimates and an assumption that all neutron stars are basically the same would tell us about their thermal evolution, which in turn would give us a hint about whether we needed exotic matter. Unfortunately, neutron stars are so small that even at the 10^6 K or higher temperatures expected for young neutron stars we can just barely detect them. Adding to the difficulty is that at those temperatures the peak emission is easily absorbed by the interstellar medium, so we can only see the high-energy tail clearly. Nonetheless, ROSAT has detected persistent X-ray emission from several young, nearby neutron stars, so now we have to interpret this emission and decide what it tells us about the star's temperature.

This ain't easy. The first complication is that the X-ray emission might not be thermal. Instead, it could be nonthermal emission from the magnetosphere. That could carry information of its own, but it makes temperature determinations difficult; basically, we have to say that, strictly, we only have upper limits on the thermal emission. Even if it were all thermal, we need to remember that we only see a section of the spectrum that is observable by an X-ray satellite, so we could be fooling ourselves about the bolometric luminosity. In fact, some early simulations of radiation transfer through a neutron star atmosphere indicated that a neutron star of effective temperature T_eff would yield far more observed counts than a blackbody at T_eff. Thus, a blackbody fit would overestimate the true temperature. These simulations used opacities computed for zero magnetic field. Thus, especially for low atomic number elements such as helium, there weren't any opacity sources at 500 eV (where the detectors operate), so in effect we would be seeing deeper into the atmosphere where it was hotter. Such simulations may be relevant for millisecond pulsars, which have magnetic fields in the 10^8 G to 10^10 G range.

Most pulsars, though, have much stronger fields, on the order of 10^12 G. In fields this strong, the binding energies of atoms go up (as mentioned before, the ground state binding energy of hydrogen in 10^12 G is 160 eV), meaning that the opacity at those higher energies rises as well. Thus, the X-ray detectors don't see as far down into the atmosphere, and the inferred temperature is less than in the nonmagnetic case. The details of the magnetic calculations are very difficult to do accurately, as they require precise computations of ionization equilibrium and polarized radiative transfer, and these are nasty in strong fields and dense, hot, matter. It seems, though, that when magnetic effects are included a blackbody isn't too bad an approximation. Stay tuned.

So what does all this mean with respect to neutron star composition? Yep, you guessed it, we don't have enough data. If you squint and look sideways at a graph of estimated temperature versus age, you might convince yourself that there is some evidence of rapid cooling, which wouldn't fit with the standard cooling scenario. But, unfortunately, the error bars are too large to be definite. We really need a large area detector that can pick up more stars.

HTH

Charles