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Part one NUCLEAR BOMBS

Sources of gravitational radiation - Let's take two stars, accelerate them to almost the speed of light and collide them. What will happen? – It will turn out to be a pretty good collider... From the forum The weakness of gravitational radiation leaves little chance for its registration. Where to look for suitable ones

So, due to the specific instability described above, large-scale gas movements occur in the convective layers of stars. The hotter masses of gas rise from the bottom up, while the colder ones fall. An intensive process of mixing the substance occurs. Calculations show, however, that the difference in the temperature of the moving elements of the gas and the environment is completely negligible, only about 1 K - and this at a temperature of the subsurface substance of the order of ten million kelvins! This is explained by the fact that convection itself tends to equalize the temperature of the layers. The average speed of rising and falling gas masses is also insignificant - only on the order of several tens of meters per second. It is useful to compare this speed with the thermal speeds of ionized hydrogen atoms in the interior of stars, which are on the order of several hundred kilometers per second. Since the speed of movement of gases participating in convection is tens of thousands of times less than the thermal speeds of particles of stellar matter, the pressure caused by convective flows is almost a billion times less than ordinary gas pressure. This means that convection does not at all affect the hydrostatic equilibrium of the stellar interior, determined by the equality of the forces of gas pressure and gravity.

You should not imagine convection as some kind of ordered process, where areas of rising gas regularly alternate with areas of its falling. The nature of convective movement is not “laminar”, but “turbulent”; that is, it is extremely chaotic, randomly changing in time and space. The chaotic nature of the movement of gas masses leads to complete mixing of the substance. This means that the chemical composition of the region of the star covered by convective movements must be homogeneous. The last circumstance is very important for many problems of stellar evolution. For example, if, as a result of nuclear reactions in the hottest (central) part of the convective zone, the chemical composition has changed (for example, there is less hydrogen, some of which has turned into helium), then in a short time this change will spread throughout the entire convective zone. Thus, “fresh” nuclear hot can continuously enter the “nuclear reaction zone” - the central region of the star - which, of course, is of decisive importance for the evolution of the star. At the same time, there may well be situations when there is no convection in the central, hottest regions of the star, which leads in the process of evolution to a radical change in the chemical composition of these regions. This will be discussed in more detail in § 12.

In § 3 we already said that the sources of energy of the Sun and stars, ensuring their luminosity during gigantic “cosmogonic” periods of time, calculated for stars of not too large a mass in billions of years, are thermonuclear reactions. Now we will dwell on this important issue in more detail.

The foundations of the theory of the internal structure of stars were laid by Eddington even when the sources of their energy were unknown. We already know that a number of important results concerning the equilibrium conditions of stars, temperature and pressure in their interiors and the dependence of luminosity on mass, chemical composition (determining the average molecular weight) and opacity of matter could be obtained without knowledge of the nature of sources of stellar energy. Nevertheless, understanding the essence of energy sources is absolutely necessary to explain the duration of the existence of stars in an almost unchanged state. Even more important is the importance of the nature of sources of stellar energy for the problem of the evolution of stars, i.e., the regular change in their main characteristics (luminosity, radius) over time. Only after the nature of the sources of stellar energy became clear, it became possible to understand the Hertzsprung-Russell diagram, the basic pattern of stellar astronomy.

The question of sources of stellar energy was raised almost immediately after the discovery of the law of conservation of energy, when it became clear that the radiation of stars is due to some kind of energy transformations and cannot continue forever. It is no coincidence that the first hypothesis about the sources of stellar energy belongs to Mayer, the man who discovered the law of conservation of energy. He believed that the source of radiation from the Sun was the continuous fall of meteoroids onto its surface. Calculations, however, showed that this source is clearly insufficient to provide the observed luminosity of the Sun. Helmholtz and Kelvin tried to explain the long-term radiation of the Sun by its slow compression, accompanied by the release of gravitational energy. This hypothesis, which is very important even (and especially!) for modern astronomy, however, turned out to be untenable to explain the radiation of the Sun over billions of years. Let us also note that at the time of Helmholtz and Kelvin there were no reasonable ideas about the age of the Sun. Only recently has it become clear that the age of the Sun and the entire planetary system is about 5 billion years.

At the turn of the 19th and 20th centuries. one of the greatest discoveries in human history was made - radioactivity was discovered. This opened up a completely new world of atomic nuclei. However, it took more than a decade for the physics of the atomic nucleus to establish a solid scientific foundation. Already by the 20s of our century it became clear that the source of energy of the Sun and stars should be sought in nuclear transformations. Eddington himself also thought so, but it was not yet possible to indicate specific nuclear processes occurring in real stellar interiors and accompanied by the release of the required amount of energy. How imperfect the knowledge of the nature of sources of stellar energy was then can be seen from the fact that Jeans, the greatest English physicist and astronomer at the beginning of our century, believed that such a source could be... radioactivity. This, of course, is also a nuclear process, but, as can easily be shown, it is completely unsuitable for explaining the radiation of the Sun and stars. This can be seen at least from the fact that such an energy source is completely independent of external conditions - after all, radioactivity, as is well known, is a process spontaneous. For this reason, such a source could not “adjust” to the changing structure of the star. In other words, there would be no “tuning” of the star's radiation. The whole picture of stellar radiation would sharply contradict observations. The first person to understand this was the remarkable Estonian astronomer E. Epic, who shortly before the Second World War came to the conclusion that only thermonuclear fusion reactions could be the source of energy for the Sun and stars.

Only in 1939 did the famous American physicist Bethe give a quantitative theory of nuclear sources of stellar energy. What kind of reactions are these? In § 7 we already mentioned that in the interior of stars there should be thermonuclear reactions. Let's look at this in a little more detail. As is known, nuclear reactions, accompanied by nuclear transformations and the release of energy, occur when particles collide. Such particles can be, first of all, the nuclei themselves. In addition, nuclear reactions can also occur when nuclei collide with neutrons. However, free (i.e., not bound in nuclei) neutrons are unstable particles. Therefore, their number in the interiors of stars should be negligible. On the other hand, since hydrogen is the most abundant element in the interior of stars and it is completely ionized, collisions of nuclei with protons will occur especially often.

In order for a proton to be able to penetrate into the nucleus with which it collides during such a collision, it needs to approach the latter at a distance of about 10 -13 cm. It is at this distance that specific attractive forces act, “cementing” the nucleus and attaching the “alien” to it. , colliding proton. But in order to approach the nucleus at such a short distance, the proton must overcome a very significant force of electrostatic repulsion (“Coulomb barrier”). After all, the nucleus is also positively charged! It is easy to calculate that to overcome this electrostatic force, the proton needs to have a kinetic energy that exceeds the potential energy of the electrostatic interaction

Stars are perhaps the most interesting thing in astronomy. In addition, their internal structure and we understand evolution better than anything in the cosmos (or so we think). The situation with planets is not very good, because their interiors are very difficult to explore - we only see what is on the surface. As for the stars, most of us are sure that they have a simple structure.

At the beginning of the last century, one young astrophysicist spoke at Eddington’s seminar in the spirit that simpler than the stars there is nothing. To which the more experienced astrophysicist replied: “Well, yes, if you are viewed from a distance of billions of kilometers, then you will also seem simple.”

In fact, stars are not as simple as they seem. But still, their properties have been studied most fully. There are two reasons for this. First, we can numerically model stars because we think they are made of ideal gas. More precisely, from plasma, which behaves like an ideal gas, the equation of state of which is quite simple. This won't work with planets. Secondly, sometimes we manage to look into the depths of stars, although so far this mainly concerns the Sun.

Fortunately, in our country there were and still are many good astrophysicists and stellar specialists. This is mainly due to the fact that there were good physicists who made nuclear weapons, and the stars are natural nuclear reactors. And when the weapon was made, many physicists, including Siberian ones, switched to studying stars, because the objects are somewhat similar. And they have written good books on this topic.

I will recommend you two books, which to this day, in my opinion, remain the best of those in Russian. “Physics of Stars,” the author of which is the famous physicist and talented teacher Samuil Aronovich Kaplan, was written almost forty years ago, but the basics have not changed since then. And modern information about the physics of stars is in the book “Stars” from the “Astronomy and Astrophysics” series, which my colleagues and I made. It enjoys such interest among readers that it has already been published in three editions. There are other books, but these two contain almost complete information for those who are becoming familiar with the subject.

Such different stars

If we look at the starry sky, we will notice that the stars have different brightness (visible brightness) and different colors. It is clear that brightness can be a matter of chance, since one star is closer, the other is further away, it is difficult to say from it what the star really is. But color tells us a lot, because the higher the body temperature, the further into the blue region the maximum in the radiation spectrum shifts. It would seem that we can simply estimate the temperature of a star by eye: red is cold, blue is hot. As a rule, this is indeed the case. But sometimes errors arise due to the fact that there is some kind of medium between the star and us. Sometimes it is very transparent, and sometimes not so much. Everyone knows the example of the Sun: high above the horizon it is white (we call it yellow, but to the eye it is almost white, because its light blinds us), but the Sun turns red when it rises or sets below the horizon. Obviously, it is not the Sun itself that changes its surface temperature, but the environment that changes its visible color, and this must be remembered. Unfortunately for astronomers this a big problem– guess how much the color has changed, i.e. the visible (color) temperature of a star, due to the fact that its light has passed through interstellar gas, the atmosphere of our planet and other absorbing media.

The spectrum of starlight is a much more reliable characteristic because it is difficult to distort too much. Everything we know about stars today we read in their spectra. The study of the stellar spectrum is a huge, carefully developed area of ​​astrophysics.

Interestingly, less than two hundred years ago one famous philosopher, Auguste Comte, said: “We have already learned a lot about nature, but there is something that we will never know - this is the chemical composition of stars, because their matter will never fall into our hands.” Indeed, it is unlikely that it will ever fall into our hands, but literally 15-20 years have passed and people have invented spectral analysis, thanks to which we learned almost everything about the chemical composition, at least, of the surface of stars. So never say never. On the contrary, there will always be a way to do something that you don’t believe in at first.

But before we talk about the spectrum, let's look again at the color of the star. We already know that the maximum intensity in the spectrum shifts to the blue region with increasing temperature, and this must be used. And astronomers have learned to use this, because taking a full spectrum is very expensive. Needed large telescope, long time observations to accumulate enough light at different wavelengths - and at the same time obtain results for only one star under study. And color can be measured very simply, and this can be done for many stars at the same time. And for mass statistical analysis, we simply photograph them two or three times through different filters with a wide transmission window.

Usually two filters - Blue (B) and Visual (V) - are enough to determine the surface temperature of a star to a first approximation. For example, we have three stars who have different temperatures surfaces, the color is different for everyone. If one of them is like the Sun (temperature about 6 thousand degrees), then in both images it will be approximately the same brightness. However, the light of a cooler star will be more attenuated by the B-filter, and little long-wavelength light will pass through it, so it will appear to us as a “weak” star. But with a hotter star the situation will be exactly the opposite.

But sometimes two filters are not enough. You can always make a mistake, like with the Sun on the horizon. Astronomers usually use 3 transmission windows: Visual, Blue, and the third - Ultraviolet, at the boundary of atmospheric transparency. Three photographs already tell us quite accurately the extent to which the interstellar medium weakens the light of each star, and what the star’s own surface temperature is. For the mass classification of stars, such 3-band photometry is so far the only method that has made it possible to study more than a billion stars.

Universal certification of stars

But the spectrum, of course, characterizes the star much more fully. The spectrum is the “passport” of a star because the spectral lines tell us so much. We are all accustomed to the words “spectral lines”; we can imagine what they are (slide 08 – spectra of chemical elements in the visible region). The horizontal axis is the wavelength, which is related to the frequency at which the light is emitted. But what is the origin of the shape of the lines, why do they look like straight vertical lines, and not circles, triangles or some kind of squiggles?

The spectral line is a monochromatic image of the entrance slit of the spectrograph. If I made a slot in the shape of a cross, I would get a set of crosses of different colors. In my opinion, a third-year physicist should think about such simple things. Or, as in the army, they said “line” - does it mean line? This is by no means always a line, because the spectrograph does not necessarily use an entrance slit, although, as a rule, the entrance hole is a vertical rectangular slit, which is more convenient.

There is always a dispersive element in the circuit of any spectrograph; a prism or diffraction grating. A star - a cloud of hot gas - emits a characteristic set of quanta of different frequencies. We pass them through the entrance slit and the dispersing element and obtain images of the slit in different colors, orderedly located along the wavelength.

If free atoms of chemical elements emit, the spectrum is lined. And if we take the hot filament of an incandescent lamp as a radiation source, then we get a continuous spectrum. Why is that? There are no characteristic energy levels in a metal conductor; electrons there, moving wildly, radiate at all frequencies. Therefore, there are so many spectral lines that they overlap each other and a continuum is obtained - a continuous spectrum.

But now we take a source of a continuous spectrum and pass its light through a cloud of gas, but colder than the spiral. In this case, the cloud snatches from the continuous spectrum those photons whose energy corresponds to transitions between energy levels in the atoms of this gas. And at these frequencies we get cut out lines, “holes” in the continuous spectrum - we get an absorption spectrum. But the atoms that absorbed light quanta became less stable and sooner or later emitted them. Why does the spectrum continue to remain “leaky”?

Because the atom doesn’t care where to throw out the “extra” energy. Spontaneous emission occurs in different directions. A certain fraction of photons, of course, flies forward, but, unlike the stimulated emission of a laser, it is tiny.

Spectral lines are usually very wide and the distribution of brightness within them is uneven. We also need to pay attention to this phenomenon and investigate what it is connected with.

There are many physical factors that make the spectral line broad. In a graph of the brightness (or absorption) distribution, two parameters can, as a rule, be distinguished: the central maximum and the characteristic width. The width of the spectral line is usually measured at the level of half the intensity of the maximum. Both the width and shape of a line can tell us about some physical features light source. But which ones?

Suppose we suspended a single atom in a vacuum and do not touch it in any way, do not prevent it from emitting. But even in this case, the spectrum will have a non-zero line width; it is called natural. It arises due to the fact that the radiation process is limited in time, for different atoms from 10⁻⁸ to 10⁻¹⁰ s. If you “cut” a sinusoid of an electromagnetic wave at the ends, then it will no longer be a sinusoid, but a curve that expands into a set of sinusoids with a continuous spectrum of frequencies. And the shorter the radiation time, the wider the spectral line.

There are other effects in natural light sources that broaden the spectral line. For example, the thermal movement of atoms. Since the radiating object has a non-zero absolute temperature, its atoms move chaotically: half towards us, half away from us, if you look at the radial velocity projection. As a result of the Doppler effect, the radiation of the first is shifted to the blue side, while the radiation of others is shifted to the red side. This phenomenon is called Doppler thermal broadening of the spectral line.

Doppler broadening can also occur for other reasons. For example, as a result of macroscopic movement of matter. The surface of any star boils: convective flows of hot gas rise from the depths, and cooled gas descends. At the moment the spectrum is taken, some flows move towards us, others - away from us. The convective Doppler effect is sometimes stronger than the thermal one.

When we look at a photograph of the starry sky, it is difficult for us to understand what the size of the stars actually is. For example, there is red and blue. If I didn't know anything about them, I might think this: a red star doesn't have a very high surface temperature, but if I see it quite bright, it means it's close to me. But then I will have a problem determining the relative distance to the blue star, which shines weaker. I think: so, blue means hot, but I don’t understand whether she is close or far from me. After all, she may be big size and emit great power, but be so far away that little light comes from there. Or, on the contrary, it can glow so faintly because it is very small, although close. How can you tell a big star from a small star? Is it possible to determine its linear size from the spectrum of a star?

It would seem not. But nevertheless it is possible! The fact is that small stars are dense, while large stars have a rarefied atmosphere, so the gas in their atmospheres is in different conditions. When we obtain the spectra of so-called dwarf stars and giant stars, we immediately see differences in the nature of the spectral lines (slide 16 - The spectra of dwarf and giant stars differ in the width of the spectral lines). In the rarefied atmosphere of the giant, each atom flies freely, rarely meeting its neighbors. They all emit almost the same way, since they do not interfere with each other, so the spectral lines of the giants have a width close to natural. But a dwarf is a massive star, but very small and, therefore, with a very high gas density. In its atmosphere, atoms constantly interact with each other, preventing their neighbors from emitting at a strictly defined frequency: because each has its own electric field, which affects the neighbor’s field. Due to the fact that the atoms are in different environmental conditions, the so-called Stark line broadening occurs. Those. By the shape, as they say, of the “wings” of the spectral lines, we immediately guess the density of the gas on the surface of the star and its typical size.

The Doppler effect can also manifest itself due to the rotation of the star as a whole. We cannot distinguish the edges of a distant star; it looks like a point to us. But from the edge approaching us, all lines of the spectrum experience a blue shift, and from the edge moving away from us, they experience a red shift (slide 18 - The rotation of a star leads to a broadening of the spectral lines). Adding up, this leads to a broadening of the spectral line. It looks different from the Stark effect and changes the shape of the spectral line differently, so you can guess in which case the line width was affected by the rotation of the star, and in which by the gas density in the star’s atmosphere. In fact, this is the only way to measure the speed of rotation of a star, because we do not see stars in the form of balls, they are all points for us.

The movement of a star in space also affects the spectrum due to the Doppler effect. If two stars move around each other, both spectra from this pair mix and appear against each other. Those. The periodic shift of lines back and forth is a sign of the orbital movement of stars.

What can we get from a series of time-varying spectra? We measure the speed (by the amplitude of the displacement), the orbital period, and from these two parameters, using Kepler’s third law, we calculate the total mass of stars. Sometimes, based on indirect evidence, it is possible to divide this mass between the components of the binary system. In most cases, this is the only way to measure the mass of stars.

By the way, the range of masses of the stars that we have studied to date is not very large: the difference is slightly more than 3 orders of magnitude. The least massive stars are about a tenth of the mass of the Sun. Their even smaller mass prevents them from triggering thermonuclear reactions. The most massive stars we have recently discovered are 150 solar masses. These are unique, only 2 of them are known so far out of several billion.

By observing rare binary systems in whose orbital plane we are located, we can also learn a lot about this pair of stars using only observational characteristics, i.e. which we can directly see, and not calculate on the basis of some laws. Since we do not distinguish them individually, we simply see a source of light, the brightness of which changes from time to time: eclipses occur while one star passes in front of the other. A deeper eclipse means that a cold star covered a hot one, and a shallower one means, on the contrary, a hot one covered a cold one (the covered areas are the same, so the depth of the eclipse depends only on their temperature). In addition to the orbital period, we measure the luminosity of stars, from which we determine their relative temperature, and from the duration of the eclipse we calculate their size.

The size of stars, as we know, is enormous. Compared to the planets, they are simply gigantic. The Sun is the most typical in size among the stars, on a par with such long-known ones as Alpha Centauri and Sirius. But the sizes of stars (as opposed to their masses) fall within a huge range - 7 orders of magnitude. There are stars noticeably smaller than them, one of the smallest (and at the same time one of the closest to us) is Proxima, it is slightly larger than Jupiter. And there are much larger stars, and at some stages of evolution they swell to incredible sizes and become noticeably larger than our entire planetary system.

Perhaps the only star whose diameter we measured directly (due to the fact that it is not far from us) is the supergiant Betelgeuse in the constellation Orion; in the Hubble telescope images it is not a dot, but a circle (slide 26 - The size of the star Betelgeuse in comparison with the diameters of the Earth and Jupiter orbits. Photo from the Hubble Space Telescope). If this star is placed in the place of the Sun, it will “eat” not only the Earth, but also Jupiter, completely covering its orbit.

But what do we even call the size of a star? Between what points do we measure the star? In optical images, the star is clearly limited in space, and it seems as if there is nothing around. So, you photographed Betelgeuse in visible light, applied a ruler to the image, and you were done? But this, it turns out, is not all. In the far infrared radiation range, it is clear that the star’s atmosphere extends much further and emits streams. We must assume that this is the boundary of the star? But we move to the microwave range and we see that the star’s atmosphere extends over almost a thousand astronomical units, several times larger than our entire Solar System.

In the general case, a star is a gaseous formation that is not enclosed in rigid walls (there are none in space) and therefore has no boundaries. Formally, any star extends indefinitely (more precisely, until it reaches a neighboring star), intensely emitting gas, which is called the stellar wind (by analogy with solar wind). Therefore, when talking about the size of a star, we always need to clarify in what range of radiation we define it, then it will be more clear what we are talking about.

Harvard Spectral Classification

The actual spectra of stars are undoubtedly very complex. They are not at all similar to the spectra of individual chemical elements that we are used to seeing in reference books. For example, even in the narrow optical range of the solar spectrum - from the violet region to the red, which our eye just sees - there are a lot of lines, and it is not at all easy to understand them. Find out, even on the basis of a detailed, highly dispersed spectrum, which chemical elements and how many stars are present in the atmosphere is a big problem that astronomers cannot fully solve.

Looking at the spectrum, we will immediately see prominent Balmer lines of hydrogen (Hα, Hβ, Hγ, Hδ) and a lot of iron lines. Sometimes helium and calcium come across. It is logical to conclude that the star consists mainly of iron (Fe) and partly hydrogen (H). At the beginning of the 20th century, radioactivity was discovered, and when people thought about the sources of energy in stars, they remembered that there are many lines of metals in the spectrum of the Sun, and assumed that the decay of uranium or radium warms the insides of our Sun. However, it turned out that this was not the case.

The first classification of stellar spectra was created at the Harvard Observatory (USA) by the hands of about a dozen women. By the way, why women in particular is an interesting question. Processing the spectra is a very delicate and painstaking job, for which the director of the observatory, E. Pickering, had to hire assistants. Women's work in science was not very welcome at that time and was paid much worse than men's: with the money that this small observatory had, it was possible to hire either two men or a dozen women. And then for the first time I was called to astronomy a large number of women who formed the so-called “Pickering harem”. The spectral classification they created was the first contribution to science by the women's team, which turned out to be much more effective than expected.

At that time, people had no idea on the basis of what physical phenomena the spectrum was formed; they simply photographed it. Trying to build a classification, astronomers reasoned like this: in the spectrum of any star there are hydrogen lines; in descending order of their intensity, all spectra can be ordered and grouped. We laid them out, designating the groups of spectra in Latin letters in alphabetical order: with the strongest lines - class A, weaker lines - class B, etc.

It seems like everything was done correctly. But a few years later quantum mechanics was born, and we realized that an abundant element is not necessarily represented in the spectrum by powerful lines, and a rare element does not manifest itself in any way in the spectrum. Much depends on the temperature.

Let's look at the absorption spectrum atomic hydrogen: only lines of the Balmer series fall into the optical range. But under what conditions are these quanta absorbed? When moving only from the second level up. But in a normal (cold) state, all the electrons “sit” in the first level, and there is almost nothing in the second. This means that we need to heat hydrogen so that some fraction of electrons jump to the second level (then they will return down again, but before that they will spend some time there) - and then the flying optical quantum can be absorbed by an electron from the second level, which will manifest itself in the visible spectrum.

So, cold hydrogen will not give us the Balmer series, but warm hydrogen will. What if we heat the hydrogen even more? Then many electrons will jump to the third and higher levels, and the second level will become depleted again. Very hot hydrogen will also not give us spectral lines that we can see in the optical range. If we go from the coldest stars to the hottest, we will see that the lines of any element can be fairly well represented in the spectrum only in a narrow temperature range.

When astrophysicists realized this, they had to rearrange the spectral types in order of increasing temperature: from cold stars to hot ones. This classification, according to tradition, is also called Harvard, but it is already natural, physical. Stars of spectral class A have a surface temperature of about 10 thousand degrees, hydrogen lines are as bright as possible, and with increasing temperature they begin to disappear, because the hydrogen atom is ionized at temperatures above 20 thousand degrees. The situation is similar with other chemical elements. By the way, in the spectra of stars colder than 4000 K there are not only lines of individual chemical elements, but also bands corresponding to molecules that are stable at such temperatures complex substances(for example, titanium and iron oxides).

The resulting sequence of letters OBAFGKM, when ordering classes by temperature, is quite easy for astronomy students to remember, especially since all sorts of mnemonic sayings have been invented. The most famous in English is Oh, Be A Fine Girl, Kiss Me! The range of surface temperatures is as follows: the hottest stars have tens of thousands of degrees, the coldest ones have a little over two thousand. For a more subtle classification, each class was divided into ten subclasses and one number from 0 to 9 was assigned to each letter on the right. Note that optical spectra in color are photographed only for beauty, but for scientific research this is meaningless, so black and white images are usually taken.

It is rare, but it happens that stars do not show absorption lines (dark on a bright background), but emission lines (bright on a dark background). Their origin is no longer so easy to understand, although this is also quite elementary. At the beginning of the lecture, we saw that a rarefied cloud of hot gas gives us emission lines. When we look at a star with emission lines in its spectrum, we understand that the source of these lines is a rarefied, translucent gas located on the periphery of the star, in its atmosphere. That is, these are stars with an extended hot atmosphere, which is transparent in the continuum (in the spaces between the lines), which means that it emits almost nothing in it (Kirchhoff’s law). But it is not transparent in individual spectral lines, and since it is not transparent in them, it emits strongly in them.

Today, the Harvard classification of stellar spectra has been expanded. New classes have been added to it, corresponding to hot stars with an extended atmosphere, the cores of planetary nebulae and novae, as well as recently discovered rather cold objects occupying intermediate position between normal stars and largest planets; they are called “brown dwarfs” or “brown dwarfs”.

There are also branches from some classes for stars with an original chemical composition. This, by the way, is a mystery to us: it is still not clear why some stars suddenly have an excess of some rare chemical element. Indeed, despite the diversity of stellar spectra, the chemical composition of their atmospheres is very similar: 98% of the mass of the Sun and similar stars consists of the first two chemical elements - hydrogen and helium, and all other elements are represented by only the remaining two percent of the mass.

The sun is the brightest source of light for us; we can stretch its spectrum very much, distinguish tens of thousands of spectral lines in it and decipher them. Thus, it has been established that all the elements of the periodic table are present on the Sun. However, I’ll tell you a secret, so far about 20 lines of the solar spectrum, very weak, have remained unidentified. So, even with the Sun, the problem of recognizing the chemical composition has not yet been completely resolved.

The distribution of chemical elements in the solar atmosphere has a number of interesting patterns). This is believed to be the typical composition of stellar matter. And for most stars this is true. Starting from carbon and up to the heaviest nuclei (at least up to uranium), there is a fairly smooth decline in the abundance of elements as their atomic number increases. However, there is a very strong gap between helium and carbon - this happens because lithium and beryllium are the easiest to participate in thermonuclear reactions, they are more active even than hydrogen and helium. And as soon as the temperature rises above a million degrees, they burn out very quickly.

But within this even trend there are peculiarities. Firstly, the peak of iron stands out sharply. In nature, including in stars, iron, nickel and elements close to them are unusually abundant compared to their neighbors. The fact is that iron is an unusual chemical element: it is the final product of thermonuclear reactions occurring under equilibrium conditions, i.e. without any explosions. In thermonuclear reactions, the star synthesizes heavier and heavier elements from hydrogen, but when it comes to iron, everything stops. Further, if we try to make something new out of iron in a thermonuclear reaction, adding neutrons, protons, and other nuclei to it, then there will be no release of heat: when the fire burns out, you won’t get anything from the ash. On the contrary, energy would have to be supplied from the outside to carry out the reaction, and no reaction with iron would take place on its own under normal conditions. Therefore, a lot of iron has accumulated in nature.

Another important point to note is that the line connecting the points on the graph has a sawtooth shape. This happens because nuclei with an even number of nucleons (protons and neutrons) are much more stable than those with an odd number. Since stable nuclei are easier to create than to destroy, there are always more of these nuclei in comparison with neighboring elements. whole order, or even one and a half.

Unlike the Sun, it contains globe and Earth-like planets contain very little hydrogen and helium, but starting with carbon, a “stellar” distribution of chemical elements is characteristic of them too. Therefore, every planet, not only the Earth, has a large iron core.

Unfortunately, the spectra only show us the composition of the surface of stars. By observing the light of a star, we can tell almost nothing about what is inside it, and the internal life of stars of different masses differs. Energy transfer in a star occurs through several mechanisms, mainly radiation and convection. For example, in stars like the Sun in the central part, where thermonuclear reactions take place, energy is mainly transferred by radiation, and the core matter does not mix with the overlying layers. Mixing occurs at the periphery, but it does not reach those internal regions in which the chemical composition gradually changes due to thermonuclear reactions. Those. thermonuclear reaction products are not carried to the surface, they circulate here starting material, from which the Sun was once born. In more massive stars, convective mixing occurs inside, but does not spread further. The accumulated chemical elements cannot jump to the surface of the star either.

Finally, low-mass stars are the most correct stars: convection is the main mechanism of heat transfer, and complete mixing of matter occurs inside them. This means that, it would seem, what has been produced in thermonuclear reactions in the center should float to their surface. However, thermonuclear reactions occur very slowly in these small stars, they spend their energy very economically and evolve slowly. Their lifespan is hundreds and thousands of times longer than that of stars like the Sun, i.e. trillions of years. And in the 14 billion years that have passed since the birth of the Universe, practically nothing has changed in their composition. They are still babies, many of them are still immature and have not started the normal thermonuclear cycle.

Thus, we still do not know what is inside the stars, what the chemical composition of the matter is there; we do not have field data. Only modeling can tell us something about this.

Hertzsprung–Russell diagram

The apparent brightness of stars is measured on an inverse logarithmic scale. magnitudes(slide 43), but for a physicist this is not interesting. What is important to him is the total radiation power of the star, and we cannot just guess it from a photograph.

For example, Alpha Centauri has amazing brightness among other stars, but this does not mean that it is the most powerful, nothing like that. This is a completely ordinary star like the Sun, it’s just that by chance it turned out to be much closer to us than the others and therefore, like a lantern, it floods the surrounding piece of sky with its light, although most of the stars neighboring it in this photo are much more powerful sources of radiation, but they are located further away.

So, we need to estimate the power of the star as accurately as possible. To do this, we use the photometric inverse square law: by measuring the apparent brightness of a star (the density of the light flux reaching the Earth) and its distance, we calculate the total power of its radiation in watts. Now we can present the general physical picture by depicting all the stars on a two-dimensional diagram (slide 46), on the axes of which two values ​​derived from observations are plotted - the temperature of the surface of the star and the relative power of its radiation (astronomers, taking into account only the optical range, call this power luminosity and measured in units of solar power). At the beginning of the 20th century, such a picture was first constructed by two astronomers, after whose names it is called the Hertzsprung–Russell diagram.

The Sun, a star with a temperature of about 6000 K and a unit power, is located almost in the middle of this diagram. Along the range of changes in both parameters, the stars are distributed almost continuously, but along the plane of the diagram they are not randomly scattered, but are grouped into compact areas.

Today, on the Hertzsprung–Russell diagram, several typical groups are distinguished, in which stars observed in nature are concentrated (slide 47). The vast majority of stars (90%) lie in a narrow band along the diagonal of the diagram; this group is called the main sequence. It ranges from dim, cool stars to hot, bright ones: from parts per million to several million solar luminosities. For a physicist, this is natural: the hotter the surface, the stronger it emits.

On both sides of the main sequence there are groups of anomalous stars. A number of high-temperature stars have unusually low luminosity (hundreds and thousands of times less than the Sun) due to their small size - we call them white dwarfs because of their color. Other exceptional stars opposite corner diagrams, is characterized by a lower temperature, but enormous luminosity - which means that they clearly have a larger physical size, these are giants.

During its evolution, a star can change its position on the diagram. More on this in one of the following lectures.