Stars are gas planets. On a clear, moonless night and in an area without light pollution, the average person can see more than 6,000 stars with the naked eye. With the help of a telescope, one can see hundreds of thousands or even millions of stars. It is estimated that there are about 150-400 billion stars in the Milky Way, and the main star of our solar system, the sun, is one star.
Two important characteristics of stars are temperature and absolute magnitude. About 100 years ago, Einar Hertzsprung of Denmark and Henry Norris Russell of the United States each drew graphs to find out whether there was a relationship between temperature and brightness. A relationship diagram is called a Hertz-Rubber diagram, or H-R diagram. In the H-R diagram, most stars form a diagonal region called the main sequence in astronomy; in the main sequence, as the absolute magnitude of a star increases, its surface temperature also increases. More than 90% of stars belong to the main sequence, and the sun is also one of these main sequence stars. Giant stars and supergiants are located higher and farther on the right side of the H-R diagram; although the surface temperature of white dwarfs is high, their brightness is not large, so they are only located in the middle and lower part of the diagram.
Stellar evolution is the continuous change of a star during its lifetime (the period of light and heat). The life span varies according to the size of the star. The evolution of a single star cannot be observed in its entirety because the processes may be too slow to be detected. So astronomers use observations of many stars at different stages of their lives and use computer models to simulate their evolution.
The astronomer Hertzsprung and the philosopher Russell first proposed the relationship between star classification and color and luminosity, established the stellar evolution relationship known as the "Hertz-Router diagram", and revealed the evolution of stars. secret. In the "H-Ro diagram", from the high temperature and strong luminosity area on the upper left to the low temperature and weak luminosity area on the lower right, there is a narrow star-dense area, including our sun; this sequence is called the main sequence , more than 90% of stars are concentrated in the main sequence. Above the main sequence region are the giant and supergiant regions; at the lower left is the white dwarf region.
By observing a star's spectrum, luminosity, and motion in space, astronomers can measure a star's mass, age, metallicity, and many other properties. The total mass of a star is the main factor that determines its evolution and final fate. Other characteristics, including diameter, rotation, motion and temperature, can be measured over the course of evolution. A diagram describing the relationship between temperature and luminosity of many stars, known as a Hertz-Rubber diagram (HR diagram), can be used to measure a star's age and stage of evolution.
Stars are not evenly distributed in galaxies. Most stars are affected by each other's gravity to form clusters of stars, such as binary stars, triple stars, or even star clusters, composed of tens of thousands to millions of stars. of stellar groups. When two binary stars orbit very close to each other, their gravitational effects may have a significant impact on their evolution. For example, a white dwarf acquires accretion disk gas from its companion star and becomes a nova. When the universe develops to a certain period, the universe is filled with uniform neutral atomic gas clouds. Large gas clouds become unstable due to their own gravity and collapse. In this way the star enters the formation stage. At the beginning of the collapse, the internal pressure of the gas cloud is very small, and the material accelerates to fall toward the center under the action of self-gravity. When the linear dimension of matter shrinks by several orders of magnitude, the situation becomes different. On the one hand, the density of the gas increases dramatically. On the other hand, due to the partial conversion of the lost gravitational potential energy into heat energy, the temperature of the gas also increases significantly. With a large increase, the pressure of a gas is proportional to the product of its density and temperature. Therefore, during the collapse process, the pressure increases faster. In this way, a pressure field sufficient to compete with the self-gravity is quickly formed inside the gas. This pressure The field finally stops the gravitational collapse, thereby establishing a new mechanical equilibrium configuration, called a star blank.
If the temperature is not high enough to ignite the hydrogen core, a brown dwarf will form.
The mechanical balance of the star base is caused by the internal pressure gradient competing with self-gravity, and the existence of the pressure gradient depends on the unevenness of the internal temperature (that is, the temperature in the center of the star base is higher than that of the periphery temperature), so thermally, this is an unbalanced system, and heat will gradually flow out from the center. This natural tendency toward thermal equilibrium has a weakening effect on mechanics.
Therefore, the star base must slowly shrink, and its gravitational potential energy decreases to increase the temperature, thereby restoring the mechanical balance; at the same time, the gravitational potential energy decreases to provide the energy required for star base radiation. This is the main physical mechanism of star blank evolution. Below we use classical gravity theory to roughly discuss this process. Consider a spherical gas cloud system with density ρ, temperature T, and radius r. The thermal motion energy of the gas is: ET= RT= T
(1) Treat the gas as a single Atomic ideal gas, μ is the molar mass, R is the gas universal constant
In order to obtain the gravitational energy Eg of the gas cloud ball, imagine that the mass of the warp ball is moved to infinity little by little, and all the balls are removed The work done by the field force is equal to -Eg. When the mass of the ball is m and the radius is r, the work done by the field force during removing dm from the surface is:
dW=- =-G( )1/3m2/3dm
( 2) So: -Eg=- ( )1/3m2/3dm= G( M5/3
So: Eg=- (2),
The total energy of the gas cloud: E=ET+EG (3) Thermal motion makes the gas distribute evenly, and gravity makes the gas concentrate. When E>0, the thermal motion dominates, and the gas cloud is stable, and small disturbances will not affect it. The gas cloud is in equilibrium; when E<0, gravity dominates, and small density disturbances cause deviations from uniformity. The gravity increases where the density is high, which intensifies the deviation and destroys the balance. The gas begins to collapse, which is obtained by E≤0. The critical radius of:
(4) The corresponding critical mass of the gas cloud is:
(5) The density of the original gas cloud is small and the critical mass is very large, so there are few stars alone. Most of them are produced by a group of stars together to form a star cluster. A spherical star cluster can contain 10^5→10^7 stars, which can be considered as being produced at the same time.
We know that the mass of the sun is MΘ=2. ×10^33, radius R=7×10^10, we put it into (2) and we can get the gravitational energy released by the sun shrinking to its current state
The total luminosity of the sun L=4×10 ^33erg.s-1 If this radiation luminosity is maintained by gravity as the energy source, then the duration is: 11×10^9 years
Many proofs show that the sun has stably maintained its current state. 5×10^9 years, therefore, the star embryo stage can only be a short transitional stage before the sun reaches a stable state like today. This raises a new question, how does the gravitational contraction of the star embryo stop? What is the energy source? The density increases during the contraction process of the main sequence star stage. We know that ρ∝r-3, from formula (4), rc∝r3/2, so rc decreases faster than r, and the shrinkage gas cloud Part of it reaches criticality under new conditions, and small disturbances can cause new local collapse. If this continues, under certain conditions, a large gas cloud shrinks into a condensate and becomes a protostar. The protostar continues to shrink after absorbing the surrounding gas clouds. The surface temperature remains unchanged and the core temperature continues to increase, causing various nuclear reactions in temperature, density and gas composition to generate heat and causing the temperature to rise extremely high. The gas pressure resists gravity to stabilize the protostar and become a star. The evolution of the star is from the main body. At the beginning of the sequence star, most of the components of the star are H and He. When the temperature reaches above 104K, that is, the average thermal kinetic energy of the particles reaches above 1eV, the hydrogen atoms are fully ionized through thermal collision (the ionization energy of hydrogen is 13.6eV) , after the temperature further increases, the collision of hydrogen nuclei in the plasma gas may cause a nuclear reaction. For high-temperature gases of pure hydrogen, the most effective nuclear reaction series is the so-called P-P chain:
The main one is. It is a 2D(p, γ)3He reaction. The content of D (deuterium, an isotope of hydrogen, consisting of one proton and one neutron) is only about 10-4% of hydrogen, and it burns out quickly (its principle is similar to modern hydrogen bomb weapons) ). If there is more D than 3He (helium 3, an isotope of helium, composed of 2 protons and 1 neutron) at the beginning, the reaction produces 3H (tritium, an isotope of hydrogen, composed of 1 proton and 2 neutrons) , which will decay into helium 3), may be the main source of 3He in the early stages of stars. This 3He that reaches the surface of stars due to convection may still remain.
Li, Be, B and other light nuclei have a very low binding energy like D, and their content is only about 2×10-9K of H. When the core temperature exceeds 3×106K, they start to burn, causing (p, α ) and (p, α) react, quickly becoming 3He and 4He. When the core temperature reaches 107K and the density reaches about 105kg/m3, the generated hydrogen is converted into He in the 41H→4He process. This is mainly the p-p and CNO cycles. Containing 1H and 4He at the same time causes a p-p chain reaction, which consists of the following three branches:
p-p1 (only 1H) p-p2 (containing both 1H and 4He) p-p3
< p> Or assume that the weight ratios of 1H and 4He are equal. As the temperature increases, the reaction gradually transitions from p-p1 to p-p3.When T>1.5×107K, the process of burning H in stars can transition to being dominated by the CNO cycle.
When heavy elements C and N are mixed in stars, they can act as catalysts to change 1H into 4He. This is the CNO cycle. The CNO cycle has two branches:
Or total The reaction rate depends on the slowest 14N (p, γ) 15O, 15N (p, α) and (p, γ) reaction branch ratio is about 2500:1.
This ratio is almost independent of temperature, so one in 2500 CNO cycles is CNO-2.
During the p-p chain and CNO cycle, the net effect is that H is burned to produce He:
Of the 26.7 MeV energy released, most of it is consumed to heat and illuminate the star, becoming The main source of stars.
We mentioned earlier that the evolution of stars begins with the main sequence, so what is the main sequence? When H is steadily burned into He, the star becomes a main sequence star. It was found that 80 to 90 percent of stars are main sequence stars. Their most common feature is that hydrogen is burning in the core region. Their luminosities, radii, and surface temperatures are all different. It was later proved that: The quantitative difference between main sequence stars is mainly their mass, followed by their age and chemical composition. This process of the sun takes about tens of millions of years.
The minimum observed mass of a main sequence star is approximately 0.1M⊙. Model calculations show that when the mass is less than 0.08M⊙, the shrinkage of the star will not reach the ignition temperature of hydrogen, and thus a main sequence star cannot be formed. This shows that there is a lower mass limit for main sequence stars. The maximum observed mass of a main sequence star is on the order of tens of solar masses. Theoretically, stars with too much mass emit strong radiation and have violent internal energy processes, so their structures are more unstable. But theoretically there is no absolute upper limit to quality.
When doing statistical analysis on a certain star cluster, people found that there is an upper limit for main sequence stars. What does this mean? We know that the luminosity of main sequence stars is a function of mass. This function can be expressed piecewise by a power formula:
L∝Mν
where υ is not a constant, its value Probably between 3.5 and 4.5. A large M reflects that there is more mass available for burning in the main sequence star, while a large L reflects the fast burning. Therefore, the lifespan of the main sequence star can be approximately marked by the trademarks of M and L:
T∝M- (ν-1)
That is, the lifespan of the main sequence star decreases according to the power law as the mass increases. If the age of the entire star cluster is T, it can be calculated from the relationship between T and M A cutoff mass MT. Main-sequence stars with masses greater than MT have ended the H-burning stage in their cores and are not main-sequence stars. This is why it is observed that star clusters composed of a large number of stars of the same age have an upper limit.
Let’s discuss the reason why most of the observed stars are main sequence stars. Table 1 is based on the constant combustion stage ignition temperature (K) of 25M⊙, the core temperature (g. cm-3) and the duration. (yr)
H 4×107 4 7×106
He 2×108 6×102 5×105
C 7×108 6×105 5 ×102
Ne 1.5×109 4×106 1
O 2×109 1×107 5×10-2
Si 3.5×109 1×108 3×10-3
The total lifetime of the combustion stage is 7.5×106
The star evolution model lists the ignition temperatures of various elements and the duration of combustion.
It can be seen from the table that the nucleus with a large atomic number has a higher ignition temperature. The largest nucleus is not only difficult to ignite, but also burns more violently after ignition, so the combustion lasts for a shorter time. Table 1 25M⊙ star evolution model of this 25M⊙. The total lifespan of the burning stage of the model star is 7.5×106 years, and more than 90% of the time is in the hydrogen burning stage, that is, the main sequence stage. Statistically speaking, this suggests that the odds of finding a star on the main sequence are higher. This is the basic reason why most of the observed stars are main sequence stars. Post-main sequence evolution: Since the main component of star formation is hydrogen, and the ignition temperature of hydrogen is lower than that of other elements, the first stage of stellar evolution is always the burning stage of hydrogen, that is, the main sequence stage. During the main sequence stage, the star maintains a stable pressure distribution and surface temperature distribution inside the star, so throughout the long stage, its luminosity and surface temperature only change slightly. Next we discuss how the star will further evolve after the hydrogen in the star core is burned out.
After the star burns out the hydrogen in the core region, it goes out. At this time, the core region is mainly helium, which is a product of combustion. The material in the outer region is mainly unburned hydrogen. After the core flameout When a star loses its radiation energy, gravitational contraction is a key factor. The end of a nuclear burning phase indicates that the temperature everywhere in the star has dropped below the temperature required to cause ignition there. The gravitational contraction will increase the temperature everywhere in the star. This is actually the search for the next nuclear ignition. At the required temperature, gravitational contraction will cause an overall increase in temperature everywhere in the star. The gravitational contraction after the main sequence will first ignite not the helium in the core (its ignition temperature is too high), but the core and periphery. There is a hydrogen shell between them. After the hydrogen shell is ignited, the core area is in a high temperature state and there is still no nuclear energy, so it will continue to shrink. At this time, due to the gravitational potential energy released by the core area and the nuclear energy released by the burning hydrogen, the non-burning hydrogen layer on the periphery must expand violently, that is, the medium radiation becomes more transparent to discharge excess heat energy. Maintain thermal balance. The expansion of the hydrogen layer reduces the surface temperature of the star, so this is a process in which the luminosity increases, the radius increases, and the surface cools. This process is the transition of the star from the main sequence to the red giant. When the process proceeds to a certain extent, the hydrogen region The temperature in the center will reach the temperature of helium ignition, and then it will transition to a new stage - the helium combustion stage.
Before helium ignition occurs in the center of the star, gravity shrinks so that its density reaches the order of 103g. cm-3. At this time, the pressure of the gas is very weakly dependent on the temperature, so the energy released by the nuclear reaction It will increase the temperature, and the increase in temperature will in turn increase the nuclear reaction rate. Once ignited, it will soon burn so violently that it will explode. This method of ignition is called a "helium flash", so in terms of phenomena You will see that the star's luminosity suddenly rises to a very high level, and then drops to a very low level.
On the other hand, when gravity contracts, its density does not reach the level of 103g. cm-3. At this time, the pressure of the gas is proportional to the temperature. The increase in ignition temperature causes the pressure to increase, and the nuclear combustion zone There will be expansion, and the expansion causes the temperature to decrease, so combustion can proceed stably, so the impact of these two ignition conditions on the evolution process is different.
How do stars evolve after a "helium flash"? The release of a large amount of energy in the flash is likely to blow away all the hydrogen in the star's outer layers, leaving behind the helium core. The helium core area reduces its density due to expansion, and helium may burn normally in it in the future. The product of helium burning is carbon. After the helium flameout, the star will have a helium shell in the carbon core area. Since the remaining mass is too small, the gravitational contraction cannot reach the ignition temperature of carbon, so it ends the evolution of helium burning, and Towards thermal death.
Since gravitational collapse is related to mass, stars with different masses evolve differently.
Stars with M<0.08M⊙: Hydrogen cannot ignite, and it will die directly without a helium burning stage.
Stars with 0.08 0.35 2.25 The reactions here are: In the early stage of the nuclear reaction, when the temperature reaches the 108K level, the 13C and 17O produced by the CNO cycle can react with 4He in a new (α, n) reaction to form 16O and 20Ne. In the nuclear reaction After a long time, 20Ne (p, γ) 21Na (β+, ν) 21Na in 21Na and 14N absorb two 4He to form 22Ne, which can react (α, n) to form 24Mg and 25Mg, etc. These reactions are as The energy source is not important, but the neutrons emitted can further cause neutron nuclear reactions. 4 After the nuclear reaction is completed, when the core temperature reaches 109K, the C, O, and Ne combustion reactions begin to occur, which are mainly C-C reactions, O-O reactions, and γ, α reactions of 20Ne: Low-mass stars (such as the sun) will expand at first. We call stars in this stage (red, blue, white) The giant star will then collapse and become a white or blue dwarf, radiate and lose energy, become a red dwarf, then a black dwarf, and eventually disappear. Massive stars, stars with ≥7 solar densities (8M⊙ Once nuclear burning stops, the star must undergo gravitational contraction. This is because the pressure inside the star to maintain mechanical equilibrium is related to its temperature. Therefore, if the star is in its final equilibrium configuration, it must be in a cold equilibrium configuration, that is, its pressure is independent of its temperature. After the core H of the main sequence star is exhausted, the stage of leaving the main sequence begins its final journey. The outcome mainly depends on the quality. For stars with very small masses, the self-gravity inside the object is not important due to their small mass. The balance inside the solid is achieved by the net Coulomb attraction between positive and negative ions and the pressure between electrons. When the mass of the star becomes larger and the self-gravity cannot be ignored, then the self-gravity increases the internal density and pressure. The increase in pressure causes the material to undergo pressure ionization, which gradually becomes the solid electricity. The confinement breaks down and the transition is to plasma gas. Increase the mass, that is, increase the density. At this time, the pressure has nothing to do with the temperature, thus reaching a cold equilibrium configuration. The kinetic energy of the electrons in the plasma is large enough to cause beta decay inside the material: Here p It is the proton in the nucleus. Such a reaction will gradually turn the nucleus in the negative ion body into a neutron-rich nucleus when the density reaches 108 g. cm-3. There will be too many neutrons in the nucleus, resulting in nuclear The structure is loose. When the density exceeds 4×1011g. cm-3, neutrons begin to separate from the nucleus and become free neutrons. The self-gravity balances the pressure between neutrons. If when the mass increases, the pressure between the neutron gas can no longer resist the self-gravity of the material, and a black hole is formed. However, due to the post-evolution stages of most stars that make the mass less than its initial mass, such as stellar winds, helium flashes, supernova explosions, etc., they will It is the star that loses a large percentage of its mass. Therefore, the fate of a star cannot be judged by its initial mass. It actually depends on the process of evolution. Then we can draw this conclusion. Stars below 8→10M⊙ eventually throw away part or most of their mass and become a white dwarf. Stars above 8M⊙ will eventually become neutron stars or black holes through the gravitational collapse of the star core. That is to say, stars with a collapsed core mass ranging from 1.44 times to 5 times that of the sun will eventually become neutron stars. Stars with a mass more than 5 times that of the sun eventually become black holes. The observed stellar mass range is generally 0.1→60M⊙. Celestial objects with a mass less than 0.08M⊙ cannot reach the ignition temperature. Therefore, if it does not emit light, it cannot become a star. The central temperature of celestial bodies with a mass greater than 60M⊙ is too high and unstable, and only less than 70 have been discovered so far. Variable magnitude. Based on actual observations and spectral analysis, we can understand the basic structure of stellar atmospheres. It is generally believed that some stars have a corona-like high-temperature and low-density corona in the outermost layer. It is often associated with star winds. Some stars have been found to have a chromosphere that produces certain emission lines in their corona. The inner atmosphere absorbs the continuous radiation of the inner high-temperature gas to form absorption lines. People sometimes call this layer of atmosphere the inversion layer, and the high-temperature layer that emits a continuous spectrum is called the photosphere. In fact, the process of forming stellar light radiation shows that the layer of the photosphere is quite thick, and each layer in it has emission and absorption. The photosphere and the reversal layer cannot be completely separated. Within the photosphere of a solar-type star, there is a convective layer that is on average about one-tenth of the radius or thicker. The position of the troposphere is very different inside upper main sequence stars and lower main sequence stars. Energy transmission is mainly by radiation in the photosphere, and by convection in the troposphere. For the photosphere and troposphere, we often use models based on actual measured physical properties and chemical compositions to conduct more detailed studies. Starting from the basic assumptions of hydrostatic equilibrium and thermodynamic equilibrium, we can establish several relationship expressions to solve for the pressure, temperature, density, opacity, productivity rate and chemical composition of different regions of the star. In the center of a star, the temperature can reach millions or even hundreds of millions of degrees, depending on the star's basic parameters and evolutionary stage. There, different productivity reactions take place. It is generally believed that stars are condensed from nebulae. Stars before the main sequence are not hot enough to undergo thermonuclear reactions and can only rely on gravitational contraction to produce energy. After entering the main sequence, the core temperature reaches more than 7 million degrees, and the thermonuclear reaction of hydrogen fusion into helium begins. This process is very long and is the longest stage in a star's life. After the hydrogen burning is completed, the star shrinks internally and expands externally, evolving into a red giant with a low surface temperature and a large volume, and may pulsate. Stars whose internal temperatures rise to nearly 100 million degrees begin to undergo a helium-carbon cycle. During these evolution processes, the temperature and luminosity of stars change according to certain rules, thus forming certain tracks on the Hertz-Rubber diagram. Finally, some stars undergo supernova explosions, their gas shells fly away, and their cores are compressed into dense stars such as neutron stars and tend to "die" (see Star Formation and Evolution).