The Milky Way and Beyond (16 page)

Supergiant stars have very extended atmospheres that are probably not even approximately in hydrostatic equilibrium. The atmospheres of M-type supergiant stars appear to be slowly expanding outward. Observations of the eclipsing binary 31 Cygni show that the K-type supergiant component has an extremely inhomogeneous, extended atmosphere composed of numerous blobs and filaments. As the secondary member of this system slowly moves behind the larger star, its light shines through larger masses of the K-type star's atmosphere. If the atmosphere were in orderly layers, the lines of ionized calcium, for example, produced by absorption of the light of the B-type star by the K-type star's atmosphere, would grow stronger uniformly as the eclipse proceeds. They do not, however.

Models of the internal structure of stars—particularly their temperature, density, and pressure gradients below the surface—depend on basic principles explained in this section. It is especially important
that model calculations take account of the change in the star's structure with time as its hydrogen supply is gradually converted into helium. Fortunately, given that most stars can be said to be examples of an “ideal gas”, the relations between temperature, density, and pressure have a basic simplicity.

Several mathematical relations can be derived from basic physical laws, assuming that the gas is “ideal” and that a star has spherical symmetry; both these assumptions are met with a high degree of validity. Another common assumption is that the interior of a star is in hydrostatic equilibrium. This balance is often expressed as a simple relation between pressure gradient and density. A second relation expresses the continuity of mass—i.e., if
M
is the mass of matter within a sphere of radius
r
, the mass added, Δ
M
, when encountering an increase in distance Δ
r
through a shell of volume 4π
r
²
Δ
r
, equals the volume of the shell multiplied by the density, ρ. In symbols,

Δ
M
= 4π
r
²
ρΔ
r
.

A third relation, termed the equation of state, expresses an explicit relation between the temperature, density, and pressure of a star's internal matter. Throughout the star the matter is entirely gaseous, and, except in certain highly evolved objects, it obeys closely the perfect gas law. In such neutral gases the molecular weight is 2 for molecular hydrogen, 4 for helium, 56 for iron, and so on. In the interior of a typical star, however, the high temperatures and densities virtually guarantee that nearly all the matter is completely ionized; the gas is said to be a plasma, the fourth state of matter. Under these conditions not only are the hydrogen molecules dissociated into individual atoms, but also the atoms themselves are broken apart (ionized) into their constituent protons and electrons. Hence, the molecular weight of ionized hydrogen is the average mass of a proton and an electron—namely, ½ on the atom-mass scale noted above. By contrast, a completely ionized helium atom contributes a mass of 4 with a helium nucleus (alpha particle) plus two electrons of negligible mass; hence, its average molecular weight is 4/3. As another example, a totally ionized nickel atom contributes a nucleus of mass 58.7 plus 28 electrons; its molecular weight is then 58.7/29 = 2.02. Since stars contain a preponderance of hydrogen and helium that are completely ionized throughout the interior, the average particle mass, μ, is the (unit) mass of a proton, divided by a factor taking into account the concentrations by weight of hydrogen, helium, and heavier ions. Accordingly, the molecular weight depends critically on the star's chemical composition, particularly on the ratio of helium to hydrogen as well as on the total content of heavier matter.

If the temperature is sufficiently high, the radiation pressure,
P
_r
, must be taken
into account in addition to the perfect gas pressure,
P
_g
. The total equation of state then becomes

P
=
P
_g
+
P
_r
.

Here
P
_g
depends on temperature, density, and molecular weight, whereas
P
_r
depends on temperature and on the radiation density constant,

a
= 7.5 × 10
^−15

ergs per cubic cm per degree to the fourth power. With μ = 2 (as an upper limit) and ρ = 1.4 grams per cubic cm (the mean density of the Sun), the temperature at which the radiation pressure would equal the gas pressure can be calculated. The answer is 28 million K, much hotter than the core of the Sun. Consequently, radiation pressure may be neglected for the Sun, but it cannot be ignored for hotter, more massive stars. Radiation pressure may then set an upper limit to stellar luminosity.

Certain stars, notably white dwarfs, do not obey the perfect gas law. Instead, the pressure is almost entirely contributed by the electrons, which are said to be particulate members of a degenerate gas. If μ' is the average mass per free electron of the totally ionized gas, the pressure,
P
, and density, ρ, are such that
P
is proportional to a 5/3 power of the density divided by the average mass per free electron; i.e.,

P
= 10
¹³
(ρ/μ')
^5/3
.

The temperature does not enter at all. At still higher densities the equation of state becomes more intricate, but it can be shown that even this complicated equation of state is adequate to calculate the internal structure of the white dwarf stars. As a result, white dwarfs are probably better understood than most other celestial objects.

For normal stars such as the Sun, the energy-transport method for the interior must be known. Except in white dwarfs or in the dense cores of evolved stars, thermal conduction is unimportant because the heat conductivity is very low. One significant mode of transport is an actual flow of radiation outward through the star. Starting as gamma rays near the core, the radiation is gradually “softened” (becomes longer in wavelength) as it works its way to the surface (typically, in the Sun, over the course of about a million years) to emerge as ordinary light and heat. The rate of flow of radiation is proportional to the thermal gradient—namely, the rate of change of temperature with interior distance. Providing yet another relation of stellar structure, this equation uses the following important quantities:
a
, the radiation constant noted above;
c
, the velocity of light; ρ, the density; and κ, a measure of the opacity of the matter. The larger the value of κ, the lower the transparency of the material and the steeper the temperature fall required to push the energy outward at the required rate. The opacity, κ, can be calculated for any temperature, density, and chemical composition and is found to depend in a
complex manner largely on the two former quantities.

In the Sun's outermost (though still interior) layers and especially in certain giant stars, energy transport takes place by quite another mechanism: large-scale mass motions of gases—namely, convection. Huge volumes of gas deep within the star become heated, rise to higher layers, and mix with their surroundings, thus releasing great quantities of energy. The extraordinarily complex flow patterns cannot be followed in detail, but when convection occurs, a relatively simple mathematical relation connects density and pressure. Wherever convection does occur, it moves energy much more efficiently than radiative transport.

The most basic property of stars is that their radiant energy must derive from internal sources. Given the great length of time that stars endure (some 10 billion years in the case of the Sun), it can be shown that neither chemical nor gravitational effects could possibly yield the required energies. Instead, the cause must be nuclear events wherein lighter nuclei are fused to create heavier nuclei, an inevitable by-product being energy.

In the interior of a star, the particles move rapidly in every direction because of the high temperatures present. Every so often a proton moves close enough to a nucleus to be captured, and a nuclear reaction takes place. Only protons of extremely high energy (many times the average energy in a star such as the Sun) are capable of producing nuclear events of this kind. A minimum temperature required for fusion is roughly 10 million K. Since the energies of protons are proportional to temperature, the rate of energy production rises steeply as temperature increases.

For the Sun and other normal main-sequence stars, the source of energy lies in the conversion of hydrogen to helium. The nuclear reaction thought to occur in the Sun is called the proton-proton cycle. In this fusion reaction, two protons (
¹
H) collide to form a deuteron (a nucleus of deuterium,
²
H), with the liberation of a positron (the electron's positively charged antimatter counterpart, denoted
e
⁺
). Also emitted is a neutral particle of very small (or possibly zero) mass called a neutrino, ν. While the helium “ash” remains in the core where it was produced, the neutrino escapes from the solar interior within seconds. The positron encounters an ordinary negatively charged electron, and the two annihilate each other, with much energy being released. This annihilation energy amounts to 1.02 megaelectron volts (MeV), which accords well with Einstein's equation

E
=
mc
²

(where
m
is the mass of the two particles,
c
the velocity of light, and
E
the liberated energy).

Next, a proton collides with the deuteron to form the nucleus of a light helium atom of atomic weight 3,
³
He. A “hard”
X-ray (one of higher energy) or gamma-ray (γ) photon also is emitted. The most likely event to follow in the chain is a collision of this
³
He nucleus with a normal
⁴
He nucleus to form the nucleus of a beryllium atom of weight 7,
⁷
Be, with the emission of another gamma-ray photon. The
⁷
Be nucleus in turn captures a proton to form a boron nucleus of atomic weight 8,
⁸
B, with the liberation of yet another gamma ray.

The
⁸
B nucleus, however, is very unstable. It decays almost immediately into beryllium of atomic weight 8,
⁸
Be, with the emission of another positron and a neutrino. The nucleus itself thereafter decays into two helium nuclei,
⁴
He. These nuclear events can be represented by the following equations:

In the course of these reactions, four protons are consumed to form one helium nucleus, while two electrons perish.

The mass of four hydrogen atoms is 4 × 1.00797, or 4.03188, atomic mass units; that of a helium atom is 4.0026. Hence, 0.02928 atomic mass unit, or 0.7 percent of the original mass, has disappeared. Some of this has been carried away by the elusive neutrinos, but most of it has been converted to radiant energy. In order to keep shining at its present rate, a typical star (e.g., the Sun) needs to convert 674 million tons of hydrogen to 670 million tons of helium every second. According to the formula

E
=
mc
²
,

more than four million tons of matter literally disappear into radiation each second.

This theory provides a good understanding of solar-energy generation, although for decades it has suffered from one potential problem. For the past several decades the neutrino flux from the Sun has been measured by different experimenters, and only one-third of flux of electron neutrinos predicted by the theory have been detected. Over that time, however, the consensus has grown that the problem and its solution lie not with the astrophysical model of the Sun but with the physical nature of neutrinos themselves. In late 1990s and early 21st century, scientists collected evidence that neutrinos oscillate between the state in which they were created in the Sun and a state that is more difficult to detect when they reach Earth.

The main source of energy in hotter stars is the carbon cycle (also called the CNO cycle for carbon, nitrogen, and oxygen), in which hydrogen is transformed into helium, with carbon serving as a catalyst. The reactions proceed as follows: first, a carbon nucleus,
¹²
C, captures a proton (hydrogen nucleus),
¹
H, to form a nucleus of nitrogen,
¹³
N, a gamma-ray photon being emitted in the process; thus,

¹²
C +
¹
H →
¹³
N + γ.

The light
¹³
N nucleus is unstable, however. It emits a positron,
e
⁺
, which encounters an ordinary electron,
e
^−
, and the two annihilate one another. A neutrino also is released, and the resulting
¹³
C nucleus is stable. Eventually the
¹³
C nucleus captures another proton, forms
¹⁴
N, and emits another gamma-ray photon. In symbols the reaction is represented by the equations