What is the significance of the main sequence

No, we must bear in mind that the survey teams counted youngsters throughout the full hour day. Some survey teams worked at night, when most youngsters were at home asleep, and others worked in the late afternoon, when most youngsters were on their way home from school and more likely to be enjoying a pizza.

If the survey was truly representative, we can conclude, however, that if an average of one-third of all youngsters are found in school, then humans ages 6 to 18 years must spend about one-third of their time in school. We can do something similar for stars. The fusion of protons to helium is an excellent, long-lasting source of energy for a star because the bulk of every star consists of hydrogen atoms, whose nuclei are protons. But if all the stars on the main sequence are doing the same thing fusing hydrogen , why are they distributed along a sequence of points?

That is, why do they differ in luminosity and surface temperature which is what we are plotting on the H—R diagram? To help us understand how main-sequence stars differ, we can use one of the most important results from our studies of model stars.

Astrophysicists have been able to show that the structure of stars that are in equilibrium and derive all their energy from nuclear fusion is completely and uniquely determined by just two quantities: the total mass and the composition of the star. This fact provides an interpretation of many features of the H—R diagram. In such a cloud, all the clumps of gas and dust that become stars begin with the same chemical composition and differ from one another only in mass.

Now suppose that we compute a model of each of these stars for the time at which it becomes stable and derives its energy from nuclear reactions, but before it has time to alter its composition appreciably as a result of these reactions. The models calculated for these stars allow us to determine their luminosities, temperatures, and sizes. If we plot the results from the models—one point for each model star—on the H—R diagram, we get something that looks just like the main sequence we saw for real stars.

And here is what we find when we do this. The model stars with the largest masses are the hottest and most luminous, and they are located at the upper left of the diagram. The least-massive model stars are the coolest and least luminous, and they are placed at the lower right of the plot. The other model stars all lie along a line running diagonally across the diagram.

In other words, the main sequence turns out to be a sequence of stellar masses. This makes sense if you think about it. The most massive stars have the most gravity and can thus compress their centers to the greatest degree. This means they are the hottest inside and the best at generating energy from nuclear reactions deep within. As a result, they shine with the greatest luminosity and have the hottest surface temperatures. The stars with lowest mass, in turn, are the coolest inside and least effective in generating energy.

Thus, they are the least luminous and wind up being the coolest on the surface. Our Sun lies somewhere in the middle of these extremes as you can see in Figure 3. The characteristics of representative main-sequence stars excluding brown dwarfs, which are not true stars are listed in Table 2. This is exactly what we found earlier when we examined the mass-luminosity relation. Our models and our observations agree. What about the other stars on the H—R diagram—the giants and supergiants, and the white dwarfs?

As a star consumes its nuclear fuel, its source of energy changes, as do its chemical composition and interior structure. These changes cause the star to alter its luminosity and surface temperature so that it no longer lies on the main sequence on our diagram.

Because stars spend much less time in these later stages of their lives, we see fewer stars in those regions of the H—R diagram. We can use the H—R diagram to explore the extremes in size, luminosity, and density found among the stars. Such extreme stars are not only interesting to fans of the Guinness Book of World Records ; they can teach us a lot about how stars work. For example, we saw that the most massive main-sequence stars are the most luminous ones.

These superluminous stars, which are at the upper left of the H—R diagram, are exceedingly hot, very blue stars of spectral type O. These are the stars that would be the most conspicuous at vast distances in space. Figure 5. The cool supergiants in the upper corner of the H—R diagram are as much as 10, times as luminous as the Sun.

In addition, these stars have diameters very much larger than that of the Sun. In contrast, the very common red, cool, low-luminosity stars at the lower end of the main sequence are much smaller and more compact than the Sun. While the core collapses, the outer layers of material in the star to expand outward. At this point the star is called a red giant.

When a medium-sized star up to about 7 times the mass of the Sun reaches the red giant phase of its life, the core will have enough heat and pressure to cause helium to fuse into carbon, giving the core a brief reprieve from its collapse. Once the helium in the core is gone, the star will shed most of its mass, forming a cloud of material called a planetary nebula.

The core of the star will cool and shrink, leaving behind a small, hot ball called a white dwarf. A white dwarf doesn't collapse against gravity because of the pressure of electrons repelling each other in its core. A red giant star with more than 7 times the mass of the Sun is fated for a more spectacular ending. These high-mass stars go through some of the same steps as the medium-mass stars.

First, the outer layers swell out into a giant star, but even bigger, forming a red supergiant. Next, the core starts to shrink, becoming very hot and dense. Then, fusion of helium into carbon begins in the core. When the supply of helium runs out, the core will contract again, but since the core has more mass, it will become hot and dense enough to fuse carbon into neon. In fact, when the supply of carbon is used up, other fusion reactions occur, until the core is filled with iron atoms.

Up to this point, the fusion reactions put out energy, allowing the star to fight gravity. However, fusing iron requires an input of energy, rather than producing excess energy. With a core full of iron, the star will lose the fight against gravity. The core temperature rises to over billion degrees as the iron atoms are crushed together.

As the Sun's core temperature is about 16 million K, the CNO cycle accounts for only a minute fraction of the total energy released. The relative energy produced by each process is shown on the plot below. How do astronomers calculate such a value? A first order approximation for this value is surprisingly easy to derive. You will recall that the mass of a helium-4 nucleus is slightly less than the sum of the four separate protons needed to form it.

A proton has a mass of 1. A helium-4 nucleus has a mass of 4. From equation 6. The production of each helium nucleus releases only a small amount of energy, 10 J which does not seem a lot. We know though measurement that the Sun's luminosity is 3. To produce this amount of energy, vast numbers of helium, 3. Each second, million tons of hydrogen fuse to form million tons of helium. This means 4 million tons of matter is destroyed and converted into energy each second.

The high temperature needed for hydrogen fusion is only found in the core region of the Sun. The energy potentially available from this mass of hydrogen is roughly:.

Given that the Sun's energy output is currently 3. As it is currently about about 5 billion years old this means it is half way through its main sequence life. We have now seen how energy is produced in a star such as the Sun. How, though, does this energy escape from the star? Two processes, radiation and convection, play a vital role. The Sun's interior comprises three main regions.

High-energy gamma photons produced in the core do not escape easily from it. The high temperature plasma in the core is about ten times denser than a dense metal on Earth. A photon can only travel a centimeter or so on average in the core before interacting with and scattering from an electron or positive ion. Each of these interactions changes both the energy and travel direction of the photon.

The direction a photon travels after an interaction is random so sometimes it is reflected back into the core. Nonetheless over many successive interactions the net effect is that the photon gradually makes its way out from the core. The path it takes is called a random walk. Photons lose energy to the electrons and ions with each interaction creating a range of photon energies. This process is known as thermalisation and results in the characteristic blackbody spectrum that forms the continuum background spectrum of stars.

Interactions between ions and electrons also produce many additional photons of various energies. These also contribute to the blackbody spectrum. The electrons and nuclei formed in fusion reactions also carry kinetic energy that they can impart to other particles through interactions, raising the thermal energy of the plasma.

Neutrinos produced by the various fusion and decay reactions travel out from the core at almost the speed of light. They are effectively unimpeded by the dense matter in the core of main sequence stars. Here, convection currents are responsible for transporting energy to the surface.

Deep cells, 30, km across are responsible for supergranulation. The cells just below the photosphere are only 1, km across and are responsible for the granulation seen on the surface of the Sun as in the image below. What happens when a main sequence star runs out of hydrogen, the fuel in its core? This leads us to evolution off the main sequence which is discussed on the next page. Skip to main content. Australia Telescope National Facility. Accessibility menu. Interface Adjust the interface to make it easier to use for different conditions.

Interface Size. High contrast mode This renders the document in high contrast mode. Invert colors This renders the document as white on black. More massive stars are hotter and bluer, while less massive stars are cooler and have a reddish appearance.

The sun falls in between the spectrum, given it a more yellowish appearance. This understanding lead to the creation of a plot known as the Hertzsprung-Russell H-R diagram, a graph of stars based on their brightness and color which in turn shows their temperature. Most stars lie on a line known as the "main sequence," which runs from the top left where hot stars are brighter to the bottom right where cool stars tend to be dimmer. Eventually, a main sequence star burns through the hydrogen in its core, reaching the end of its life cycle.

At this point, it leaves the main sequence. Stars smaller than a quarter the mass of the sun collapse directly into white dwarfs. White dwarfs no longer burn fusion at their center, but they still radiate heat. Eventually, white dwarfs should cool into black dwarfs , but black dwarfs are only theoretical; the universe is not old enough for the first white dwarfs to sufficiently cool and make the transition.

Larger stars find their outer layers collapsing inward until temperatures are hot enough to fuse helium into carbon. Then the pressure of fusion provides an outward thrust that expands the star several times larger than its original size, forming a red giant. The new star is far dimmer than it was as a main sequence star.