astrtology homework 10 questions

chapter10.ppt

Home >Geography homework help >astrtology homework 10 questions

Michael Seeds Dana Backman

Chapter 10
Structure and Formation of Stars

Jim he allowed [the stars] was made, but I allowed they happened. Jim said the moon could’a laid them; well, that looked kind of reasonable, so I didn’t say nothing against it, because I’ve seen a frog lay most as many, so of course it could be done.

MARK TWAIN

The Adventures of Huckleberry Finn

The stars are not eternal.
When you look at the sky, you see hundreds of points of light, and each is an object like the sun held together by its gravity and generating tremendous energy in its core through nuclear reactions.
The stars you see tonight are the same stars your parents, grandparents, and great-grandparents saw.
Stars hardly change at all in a human lifetime, but they are not eternal.
Stars are born and stars die.
This chapter begins that story.

How can you know what the internal processes and lifecycles of stars are when you can’t see inside them and don’t live long enough to see them evolve?
The answer lies in the methods of science.
By constructing theories that describe how nature works and then testing those theories against evidence from observations, you can unravel some of nature’s greatest secrets.

In this chapter, you will see how the flow of energy from inside to the surfaces of the stars balances gravity—making the stars stable—and how nuclear reactions inside stars generate that energy.
You will also see how gravity creates new stars from the thin dusty gas between the stars.

If there is a single idea in modern astronomy that can be called critical, it is the concept of balance.
Stars are simple, elegant power sources held together by their own gravity and balanced by the support of their internal heat and pressure.

Stellar Structure

Using the basic concept of balance, you can consider the structure of a star.
This comprises the variation in temperature, density, pressure, and other such aspects between the surface of a star and its center.

Stellar Structure

It will be easier to think about stellar structure if you imagine that the star is divided into concentric shells—such as those in an onion.
You can then discuss the temperature, pressure, and density in each shell.
Of course, these helpful shells exist only in the imagination—stars have no separable layers.

Stellar Structure

The first two laws of stellar structure have some thing in common—they are both what astronomers and physicists call conservation laws.
They state that certain things cannot be created out of nothing or vanish into nothing.
Such conservation laws are powerful aids to help you understand how nature works.

The Laws of Conservation of Mass and Energy

The law of conservation of mass is a basic law of nature that can be applied to the structure of stars.
It states that the total mass of a star must equal the sum of the masses of its shells.
This is like saying the weight of a cake must equal the sum of the weight of its layers.

The Laws of Conservation of Mass and Energy

The law of conservation of energy is another basic law of nature.
It states that the amount of energy flowing out of the top of a layer in the star must be equal to the amount of energy coming in at the bottom plus whatever energy is generated within the layer.

The Laws of Conservation of Mass and Energy

That means that the energy leaving the surface of the star—its luminosity—must equal the sum of the energies generated in all the layers inside the star.
This is like saying that all the new cars driving out of a factory must equal the sum of all the cars made on each of the production lines.

The Laws of Conservation of Mass and Energy

The two laws may seem so familiar and so obvious that they hardly need to be stated.
However, they are important clues to the structure of stars.

The Laws of Conservation of Mass and Energy

The weight of each layer must be supported by the layer below.
The words down or below are conventionally used to refer to regions closer to the center of a star.

Hydrostatic Equilibrium

Picture a pyramid of people in a circus stunt.
The people in the top row do not have to hold up anybody else.
The people in the next row down are holding up the people in the top row.

Hydrostatic Equilibrium

In a star that is stable, the deeper layers must support the weight of all the layers above.
As the inside of a star is made up of gas, the weight pressing down on a layer must be balanced by the gas pressure in the layer.
If the pressure is too low, the weight from above will compress and push down the layer.
If the pressure is too high, the layer will expand and lift the layers above.

Hydrostatic Equilibrium

This balance between weight and pressure is called hydrostatic equilibrium.
Hydro (from the Greek word for water) informs you the material is a fluid—which, by definition, includes the gases of a star.
Static informs you the fluid is stable—neither expanding nor contracting.

Hydrostatic Equilibrium

The figure shows this
hydrostatic balance in
the imaginary layers of
a star.
The weight pressing down on
each layer is shown by lighter
red arrows, which grow larger
with increasing depth because the weight grows larger.
The pressure in each layer is
shown by darker red arrows,
which must grow larger with
increasing depth to support
the weight.

Hydrostatic Equilibrium

The pressure in a gas depends on the temperature and density of the gas.
Near the surface, there is little weight pressing down—so, the pressure need not be high for stability.
Deeper in the star, the pressure must be higher—which means that the temperature and density of the gas must also be higher.
Hydrostatic equilibrium informs you that stars must have high temperature, pressure, and density inside to support their own weight and be stable.

Hydrostatic Equilibrium

Although the law of hydrostatic equilibrium can inform you of some things about the inner structure of stars, you need one more law to completely describe a star.
You need to know how energy flows from the inside to the outside.

Hydrostatic Equilibrium

The surface of a star radiates light and heat into space—and would quickly cool if that energy were not replaced.
As the inside of the star is hotter than the surface, energy must flow outward to the surface, where it radiates away.
This flow of energy through each shell determines its temperature—which, as you have learned, determines how much weight that shell can balance.

Energy Transport

To understand the structure of a star, you must understand how energy moves from the center to the surface.

Energy Transport

The law of energy transport states that energy must flow from hot regions to cooler regions by conduction, convection, or radiation.

Energy Transport

Conduction is the most familiar form of heat flow.
If you hold the bowl of a spoon in a candle flame, the handle of the spoon grows warmer.
Heat, in the form of motion
among the atoms of the
spoon, is conducted from
atom to atom up the
handle.

Energy Transport

Thus, conduction requires close contact between the atoms.
So, it is not significant in normal stars that do not have high enough densities for conduction even in their centers.

Energy Transport

The transport of energy by radiation is another familiar experience.
Put your hand near a candle flame, and you can feel the heat.
What you actually feel are infrared
photons—packets of energy—
radiated by the flame and absorbed
by your hand.

Energy Transport

Radiation is the principal means of energy transport in the interiors of most stars.
Photons are absorbed and reemitted in random directions over and over as energy works its way from the hot interior toward the cooler surface.

Energy Transport

The flow of energy by radiation depends on how difficult it is for photons to move through the gas.
The opacity of the gas—its resistance to the flow of radiation—depends strongly on its temperature.
A hot gas is more transparent than a cool gas.

Energy Transport

If the opacity is high, radiation cannot flow through the gas easily—and it backs up like water behind a dam.
When enough heat builds up, the gas begins to churn as hot gas rises upward and cool gas sinks downward.

Energy Transport

This heat-driven circulation of a fluid is convection, the third way energy can move in a star.
The rising wisp of smoke above
a candle flame is carried by convection.
Energy is carried upward in convection
currents as rising hot gas and also as
sinking cool gas.

Energy Transport

The figure shows a cross-section of the sun in which zones of radiative and convective energy transport are indicated.

Energy Transport

The four laws of stellar structure are summarized in the table.
Properly understood, they can inform you how stars are born, how they live, and how they die.

Energy Transport

The laws of stellar structure can be written as mathematical equations.
By solving the equations, astronomers can make a mathematical model of the inside of a star.

Stellar Models

To start one of these models requires knowing the type of information you learned earlier about the external properties of a star:
Diameter
Luminosity
Mass
Surface temperature
Atmospheric density
Composition

Stellar Models

If you want to build a model of a star, you can divide the star into, say, 100 concentric shells and then write down the four equations of stellar structure for each shell.
You would then have 400 equations that would have 400 unknowns, namely, the temperature, density, mass, and energy flow in each shell.

Stellar Models

Solving 400 equations simultaneously is not easy.
The first such solutions, done by hand before the invention of the electronic computer, took months of work.

Stellar Models

Now, a properly programmed computer can solve the equations in a few seconds and print a table of numbers that represents the conditions in each shell of the star.
Such a table is called a stellar model.

Stellar Models

The table in the figure is a model of the sun.
The bottom line, for radius equal to 0.00, represents the center of the sun.
The top line, for radius equal to 1.00, represents the surface.

Stellar Models

The other lines represent the temperature and density in each shell, the mass inside each shell, and the fraction of the sun’s luminosity flowing outward through the shell.

Stellar Models

The bottom line informs you that the temperature at the center of the sun must be over 15 million Kelvin in order for the sun to be stable.
At such a high temperature, the gas is highly transparent, and energy flows as radiation.

Stellar Models

Nearer the surface, the temperature is lower, the gas is more opaque, and energy moves by convection.
This is confirmed by evidence of convection visible on the sun’s surface.

Stellar Models

Stellar models also let you look into a star’s past and future.
You can use them as time machines to follow the evolution of stars over billions of years.

Stellar Models

For instance, to look into a star’s future, you could use a model to determine how fast the star uses its fuel in each shell.
As the fuel is consumed, the chemical composition of the gas changes and the amount of energy generated declines.
By calculating the rate of these changes, you could predict what the star will look like at some point in the future.

Stellar Models

You could repeat the process over and over and step-by-step follow the evolution of the star as it ages.

Stellar Models

Stellar astronomy has made great advances since the 1960s.
High-speed computers began to be available that could calculate models of the structure and evolution of stars.

Stellar Models

The summary of star formation in this chapter is based on thousands of stellar models.
You will continue to rely on theoretical models as you study main-sequence stars later in the chapter.

Stellar Models

Astronomers often use the wrong words to describe energy generation in the sun and stars.
Astronomers will say, “The star ignites hydrogen burning.”
You normally use the word ignite to mean ‘catch on fire,’ and burn means ‘consume by fire.’
What goes on inside stars isn’t really burning in the usual sense.

Nuclear Fusion in the Sun and Stars

The sun is a normal star, and it is powered by nuclear reactions that occur near its center.
The energy produced keeps the interior hot and the gas is totally ionized.
That is, the electrons are not attached to atomic nuclei.
The gas is a soup of rapidly moving particles colliding with each other at high velocity.

Nuclear Fusion in the Sun and Stars

When astronomers discuss nuclear reactions inside stars, they refer to atomic nuclei and not to atoms.

Nuclear Fusion in the Sun and Stars

How exactly can the nucleus of an atom yield energy?
The answer lies in the forces that hold the nuclei together.

Nuclear Fusion in the Sun and Stars

The sun generates its energy by breaking and reconnecting the bonds between the particles inside atomic nuclei.
This is quite different from the way you generate energy by burning wood in a fireplace.
The process of burning extracts energy by breaking and reconnecting chemical bonds between atoms in the fuel.

Nuclear Binding Energy

Chemical bonds are formed by the electrons on the outside of atoms.
You have learned that the electrons are bound to the atoms by the electromagnetic force.
Thus, chemical energy originates in the electromagnetic force.

Nuclear Binding Energy

There are only four known forces in nature—the force of gravity, the electromagnetic force, and the strong and weak nuclear forces.
The weak force is involved in the radioactive decay of certain kinds of nuclear particles.
The strong force binds together atomic nuclei.
Thus, nuclear energy originates in the strong nuclear force.

Nuclear Binding Energy

Nuclear power plants on Earth generate energy through nuclear fission reactions that split uranium nuclei into less massive fragments.
A uranium nucleus contains a total of 235 protons and neutrons.
It splits into a range of fragments containing roughly half as many particles.
As the fragments produced are more tightly bound than the uranium nuclei, binding energy is released during uranium fission.

Nuclear Binding Energy

Stars make energy in nuclear fusion reactions that combine light nuclei into heavier nuclei.
The most common reaction, including that in the sun, fuses hydrogen nuclei (single protons) into helium nuclei (two protons and two neutrons).
As the nuclei produced are more tightly bound than the original nuclei, energy is released.

Nuclear Binding Energy

The graph shows that both fusion and fission reactions move downward in the diagram.
That means they release the binding energy of nuclei and create more tightly bound nuclei.

Nuclear Binding Energy

The sun fuses four hydrogen nuclei to make one helium nucleus.
As one helium nucleus contains 0.7 percent less mass than four hydrogen nuclei, some mass vanishes in the process.
That mass is converted to energy.
You can figure out how much by using Einstein’s famous equation, E = mc2.

Hydrogen Fusion

For example, converting 1 kg (2.2 lb) of matter completely into energy would produce an enormous amount of energy—9 x 1016 Joules.
This is comparable to the amount of energy a big city like Philadelphia uses in a year.

Hydrogen Fusion

Making one helium nucleus releases only a small amount of energy—hardly enough to raise a housefly 1/40 of a millimeter into the air.
Only by concentrating many reactions in a small region can nature produce significant amounts of energy.

Hydrogen Fusion

The sun has a voracious appetite and needs 1038 reactions per second, transforming 5 million tons of mass into energy every second, just to replace the energy pouring into space from its surface.

It might sound as if the sun is losing mass at a furious rate.

However, during its entire 10-billion-year lifetime, the sun will convert less than 0.07 percent of its mass into energy.

Hydrogen Fusion

Fusion reactions can occur only when the nuclei of two atoms get very close to each other.
As atomic nuclei carry positive charges, they repel each other with an electrostatic force called the Coulomb force.
Physicists commonly refer to this repulsion between nuclei as the Coulomb barrier.

Hydrogen Fusion

To overcome this barrier, atomic nuclei must have violent collisions that are rare unless the gas is very hot, in which case the nuclei move at high speeds.

Hydrogen Fusion

So, nuclear reactions in the sun take place only near the center—where the gas is hot and dense.
A high temperature ensures that collisions between nuclei are violent enough to overcome the Coulomb barrier.
A high density ensures that there are enough collisions, and thus enough reactions, to produce enough energy to keep the sun stable.

Hydrogen Fusion

You can symbolize this process with a simple nuclear reaction:

4 1H → 4He + energy

1H represents a proton, the nucleus of a hydrogen atom.
4He represents the nucleus of a helium atom.
The superscripts indicate the total number of protons and neutrons in each nucleus.

Hydrogen Fusion

The actual steps in the process are more complicated.
Instead of waiting for four hydrogen nuclei to collide simultaneously, a highly unlikely event, the process can proceed step by step in a chain of reactions called the proton–proton chain.

Hydrogen Fusion

The proton-proton chain is a series of three nuclear reactions that builds a helium nucleus by adding protons one at a time.

Hydrogen Fusion

This process is efficient at temperatures above 10,000,000 K.
The sun, for example, manufactures its energy in this way.

Hydrogen Fusion

The three reactions in the proton–proton chain are:

1H + 1H → 2H + e+ + v

2H + 1H → 3He + γ

3He + 3He → 4He + 1H + 1H

Hydrogen Fusion

In the first reaction, two hydrogen nuclei (two protons) combine, and one changes into a neutron—to result in a heavy hydrogen nucleus called deuterium.
This emits a particle called a positron (a positively charged electron, symbolized by e+) and another called a neutrino (ν).

1H + 1H → 2H + e+ + ν

Hydrogen Fusion

In the second reaction, the heavy hydrogen nucleus absorbs another proton and, with the emission of a gamma ray (γ), becomes a lightweight helium nucleus (3He).

2H + 1H → 3He + γ

Hydrogen Fusion

Finally, two light helium nuclei combine to form a normal helium nucleus and two hydrogen nuclei.

3He + 3He → 4He + 1H + 1H

Hydrogen Fusion

As the last reaction needs two 3He nuclei, the first and second reactions must each occur twice for each 4He produced.

Hydrogen Fusion

The net result of this chain reaction is the transformation of four hydrogen nuclei into one helium nucleus plus energy.

Hydrogen Fusion

All main-sequence stars fuse hydrogen into helium to generate energy.
The sun and smaller stars fuse hydrogen by the proton–proton chain.
Upper-main-sequence stars, more massive than the sun, fuse hydrogen by a more efficient process called the CNO (carbon-nitrogen-oxygen) cycle.

Hydrogen Fusion

The CNO cycle begins with a carbon nucleus, transforms it first into a nitrogen nucleus, then into an oxygen nucleus, and then back to a carbon nucleus.
The carbon is unchanged in the end.
However, along the way, four hydrogen nuclei are fused to make a helium nucleus plus energy—just as in the proton–proton chain.

Hydrogen Fusion

A carbon nucleus has six times more positive electric charge than hydrogen.
So, the Coulomb barrier is higher than for combining two protons.

Hydrogen Fusion

Temperatures higher than 16,000,000 K are required to make the CNO cycle work.
The center of the sun is not quite hot enough for this reaction.
In stars more massive than about 1.1 solar masses, the cores are hot enough, and the CNO cycle dominates over the slower proton–proton chain.

Hydrogen Fusion

In both the proton–proton chain and the CNO cycle, energy appears in the form of gamma rays, positrons, and neutrinos.

Hydrogen Fusion

The gamma rays are photons that are absorbed by the surrounding gas before they can travel more than a few centimeters.
This heats the gas.

Hydrogen Fusion

The positrons produced in the first reaction combine with free electrons.
Both particles vanish, converting their mass into more gamma rays.
Thus, the positrons also help keep the center of the star hot.

Hydrogen Fusion

The neutrinos travel at nearly the speed of light and almost never interact with other particles.
The average neutrino could pass unhindered through a lead wall a light-year thick.
Thus, the neutrinos do not help heat the gas but race out of the star, carrying away roughly 2 percent of the energy produced.

Hydrogen Fusion

It is time to ask the critical question that lies at the heart of science.
What is the evidence to support this theoretical explanation of how the sun and other stars make energy?

Hydrogen Fusion

The search for that evidence will introduce you to one of the great triumphs of modern astronomy.

Hydrogen Fusion

Nuclear reactions in the sun’s core produce floods of neutrinos that rush off into space.
If you could detect these neutrinos, you could probe the sun’s interior.

Neutrinos from the Sun’s Core

As they almost never interact with atoms, neutrinos are extremely hard to detect.
You never feel the flood of over 1012 solar neutrinos that flows through your body every second.
Even at night, neutrinos from the sun rush through Earth as if it weren’t there, up through your bed, through you, and onward into space.

Neutrinos from the Sun’s Core

Certain nuclear reactions, however, can be triggered by a neutrino of the right energy.
Beginning in the 1960s, astronomers found several ways to detect solar neutrinos.

Neutrinos from the Sun’s Core

Initial results, however, detected too few neutrinos—about one-third as many as predicted by models of the sun’s interior.
Later experiments were finally able to solve the puzzle and confirm that the original models were correct.

Neutrinos from the Sun’s Core

Physicists know of three kinds of neutrinos, which they call ‘flavors.’
The first neutrino experiments could detect (or taste) only one flavor, electron-neutrinos.

Neutrinos from the Sun’s Core

Theory hinted that the electron-neutrinos produced in the core of the sun might change back and forth among the three flavors as they rushed out through the sun and across space to Earth.

Neutrinos from the Sun’s Core

Observations begun in 2000 confirm this.
Two-thirds of the electron
neutrinos produced in the sun
transform into muon- and
tau-neutrinos, which most
detectors cannot count,
before they reach Earth.

Neutrinos from the Sun’s Core

You can be confident that models of the sun’s interior are basically correct.
They predict the true rate of proton–proton fusion in the sun’s core.

Neutrinos from the Sun’s Core

Nuclear reactions in stars manufacture energy and heavy atoms under the supervision of a natural thermostat that keeps the reactions from erupting out of control.
That thermostat is the relation between gas pressure and temperature.

The Pressure–Temperature Thermostat

In a star, the nuclear reactions generate just enough energy to balance the inward pull of gravity.
Consider what would happen if they produced too much energy.
As the star balances gravity by generating energy, the extra energy flowing out of the star would force it to expand.
The expansion would lower the central temperature and density and slow the nuclear reactions until the star regained stability.

The Pressure–Temperature Thermostat

Thus, the star has a built-in regulator that keeps the nuclear reactions from occurring too rapidly.
The same thermostat keeps the reactions from dying down.
Suppose the nuclear reactions began making too little energy.
Then, the star would contract slightly—increasing the central temperature and density and increasing the nuclear energy generation.

The Pressure–Temperature Thermostat

The overall stability of a star depends on the relation between gas pressure and temperature.
If the material of the star has the property of normal gases—for which an increase or decrease in temperature produces a corresponding change in pressure—then the nuclear reaction pressure–temperature thermostat can function properly and contribute to the stability of the star.

The Pressure–Temperature Thermostat

You can understand the deepest secret of the stars by looking at the most obvious feature of the H-R diagram—the main sequence.

Main-Sequence Stars

Roughly 90 percent of all the stars that make energy by nuclear fusion are main-sequence stars.
They obey a simple relationship that is the key to understanding the life and death of stars.

Main-Sequence Stars

Observations of the temperature and luminosity of stars show that main-sequence stars obey a simple rule.
The more massive a star is, the more luminous it is.

The Mass–Luminosity Relation

This rule, the mass–luminosity relation, is the key to understanding the stability of main-sequence stars.
In fact, the mass–luminosity relation is predicted by theories of stellar structure, giving astronomers direct observational confirmation of those theories.

The Mass–Luminosity Relation

To understand the mass–luminosity relation, you must consider both the law of hydrostatic equilibrium, which states that pressure balances weight, and the pressure–temperature thermostat, which regulates energy production.

The Mass–Luminosity Relation

A star that is more massive than the sun has more weight pressing down on its interior.
So, the interior must have a high pressure to balance that weight.
That means the massive star’s automatic pressure–temperature thermostat must keep the gas in its interior hot and the pressure high.

The Mass–Luminosity Relation

A star less massive than the sun has less weight on its interior and thus needs less internal pressure.
Thus, its pressure–temperature thermostat is set lower.
So, massive stars are more luminous because they must support their large weight.
If they were not so luminous, they would not be stable.

The Mass–Luminosity Relation

The mass–luminosity relation also shows you why the main sequence has a lower end—a minimum mass.
Objects with masses less than about 0.08 of the sun’s mass cannot raise their central temperature high enough to begin hydrogen fusion.

Brown Dwarfs

Called brown dwarfs, such objects are only about ten times the diameter of Earth—about the same size as Jupiter—and may still be warm from the processes of formation.
However, they do not generate energy by hydrogen fusion.
They have contracted as much as they can and are slowly cooling off.

Brown Dwarfs

Brown dwarfs fall in the gap between low-mass M stars and massive planets like Jupiter.
They would look dull orange or red to your eyes—which is why they are labeled ‘brown’—because they emit most of their energy in the infrared.
The warmer brown dwarfs fall in spectral class L and the cooler in spectral class T.

Brown Dwarfs

As they are so small and cool, brown dwarfs are very low-luminosity objects and difficult to find.
Nevertheless, hundreds are known, and they may be as common as M stars.

Brown Dwarfs

While a star is on the main sequence, it is stable—so, you might think its life would be totally uneventful.
However, a main-sequence star balances its gravity by fusing hydrogen.
As the star gradually uses up its fuel, that balance must change.
Thus, even stable main-sequence stars are changing as they consume their hydrogen fuel.

The Life of a Main-Sequence Star

You have learned that hydrogen fusion combines four nuclei into one.
Thus, as a main-sequence star fuses its hydrogen, the total number of particles in its interior decreases.

The Life of a Main-Sequence Star

Each newly made helium nucleus can exert the same pressure as a hydrogen nucleus.
As the gas has fewer nuclei, though, its total pressure is less.
This unbalances the gravity–pressure stability, and gravity squeezes the core of the star more tightly.

The Life of a Main-Sequence Star

As the core contracts, its temperature increases and the nuclear reactions run faster—releasing more energy.
This additional energy flowing outward forces the outer layers to expand.
Therefore, as a main-sequence star slowly turns hydrogen into helium in its core, the core contracts and heats up while the outer parts of the star become larger, the star becomes more luminous, and its surface cools down.

The Life of a Main-Sequence Star

These gradual changes during the lifetimes of main-sequence stars mean that the main sequence is not a sharp line across the H-R diagram but rather a band.

The Life of a Main-Sequence Star

Stars begin their stable lives fusing hydrogen on the lower edge of the band, known as the zero-age main sequence (ZAMS).

The Life of a Main-Sequence Star

However, gradual changes in luminosity and surface temperature move them upward and slightly to the right.

The Life of a Main-Sequence Star

By the time they reach the upper edge of the main sequence, they have exhausted nearly all the hydrogen in their centers.

The Life of a Main-Sequence Star

If you precisely measure and plot the luminosity and temperature of main-sequence stars on the diagram, you will find them at various positions within the main-sequence band.

The Life of a Main-Sequence Star

That indicates how much hydrogen has been converted to helium in their cores.

The Life of a Main-Sequence Star

Therefore, you can use the position of a star in the band, combined with stellar evolution models, as one way to estimate the star’s age.

The Life of a Main-Sequence Star

The sun is a typical main-sequence star.
As it undergoes these gradual changes, Earth will suffer.

The Life of a Main-Sequence Star

When the sun began its main-sequence life about 5 billion years ago, it was only about 70 percent as luminous as it is now.
By the time it leaves the main sequence in another 5 billion years, it will have twice its present luminosity.
Long before that, the rising luminosity will raise Earth’s average temperature, melt the polar caps, modify Earth’s climate, and ultimately drive away the oceans and atmosphere.

The Life of a Main-Sequence Star

Life on Earth will probably not survive these changes.
Nevertheless, you have a billion years or more to prepare.

The Life of a Main-Sequence Star

The average star spends 90 percent of its life on the main sequence.
This explains why 90 percent of all true stars are main-sequence stars.
You are most likely to see a star during that long, stable period while it is on the main sequence.

The Life of a Main-Sequence Star

The number of years a star spends on the main sequence depends on its mass.
Massive stars consume fuel rapidly and live short lives.
In contrast, low-mass
stars conserve their
fuel and shine for
billions of years.

The Life of a Main-Sequence Star

For example, a 25-solar-mass star will exhaust its hydrogen and die in only about 7 million years.
The sun has enough fuel to last about 10 billion years.
The red dwarfs, although they have little fuel, use it up very slowly and may be able to survive for 100 billion years or more.

The Life of a Main-Sequence Star

The key to understanding star formation is the correlation between young stars, short-lived stars, and interstellar clouds of gas and dust.

The Birth of Stars

Where you find the youngest groups of stars as well as stars with very short lifetimes, you also find large clouds of gas illuminated by the hottest and brightest of the new stars.

The Birth of Stars

This should lead you to suspect that stars form from such clouds.

The Birth of Stars

A common misperception is to imagine that space is empty—a vacuum.

The Interstellar Medium

In fact, if you glance at Orion’s sword on a winter evening, you can see the Great Nebula in Orion—a glowing cloud of gas and dust.

The Interstellar Medium

That cloud and others show that the space between the stars is filled with low-density gas and dust called the interstellar medium (ISM).
About 75 percent of the mass of the gas is hydrogen.
25 percent is helium.
There are also traces of carbon, nitrogen, oxygen, calcium, sodium, and heavier atoms.
This is very similar to the composition of the sun and other stars.

The Interstellar Medium

Roughly 1 percent of the ISM’s mass is made up of microscopic dust called interstellar dust.
The dust particles are about the size of the particles in cigarette smoke.
The dust seems to be made mostly of carbon and silicates (rocklike minerals) mixed with or coated with frozen water plus some organic compounds.

The Interstellar Medium

This interstellar material is not uniformly distributed through space.
It consists of a complex tangle of cool, dense clouds pushed and twisted by currents of hot, low-density gas.
Although the cool clouds contain only 10 to 1,000 atoms/cm3 (fewer than in any laboratory vacuum on Earth), astronomers refer to them as dense clouds in contrast with the hot, low-density gas that fills the spaces between clouds.
That thin gas contains only about 0.1 atom/cm3.

The Interstellar Medium

How do you know there is an interstellar medium?
How do you know its properties?

The Interstellar Medium

In some cases, the interstellar medium is easily visible as clouds of gas and dust—as in the case of the Great Nebula in Orion.
Astronomers call such
a cloud a nebula—
from the Latin word
for cloud.
Such nebulae give
clear evidence of an
interstellar medium.

The Interstellar Medium

There are three kinds of nebulae.
There are three important points to note about nebulae.

The Interstellar Medium

One, very hot stars can excite clouds of gas and dust to emit light.
This reveals that the clouds contain mostly hydrogen gas at very low densities.

The Interstellar Medium

Two, where dusty clouds reflect the light of slightly cooler stars, you see evidence that the dust in the clouds is made up of very small particles.

The Interstellar Medium

Three, some dense clouds of gas and dust are detectable only where they are silhouetted against background regions filled with stars or bright nebulae.

The Interstellar Medium

If a cloud is less dense, the starlight may be able to penetrate it, and stars can be seen through the cloud.
However, the stars look dimmer because the dust in the cloud scatters some of the light.

The Interstellar Medium

As shorter wavelengths are scattered more easily than longer wavelengths, the redder photons are more likely to make it through the cloud.

The Interstellar Medium

Thus, the stars look slightly redder than they should for their respective spectral types.
This effect is called
interstellar reddening.
The same process makes
the setting sun look redder.

The Interstellar Medium

The fact that distant stars are dimmed and reddened by intervening gas and dust is clear evidence of an interstellar medium.

The Interstellar Medium

Although you have been imagining individual nebulae, the thin gas and dust of the interstellar medium also fills the spaces between the nebulae.

The Interstellar Medium

You can see evidence of that because the gas forms interstellar absorption lines in the spectra of distant stars.
As starlight travels through the interstellar medium, gas atoms of elements such as calcium and sodium absorb photons of certain wavelengths—producing absorption lines.

The Interstellar Medium

You can be sure they originate in the interstellar medium because they appear in the spectra of O and B stars.
These stars are too hot to have visible calcium and sodium absorption lines in their own atmospheres.

The Interstellar Medium

Also, the narrowness of the interstellar lines indicates they could not have been formed in the hot atmospheres of the stars.
The widths indicate a very small range of Doppler shifts.
This means the
atoms are moving
at speeds
corresponding to
temperatures of
only 10 to 50 K.

The Interstellar Medium

Observations at nonvisible wavelengths provide valuable evidence about the interstellar medium.

The Interstellar Medium

Infrared observations can directly detect dust in the interstellar medium.
Although the dust grains are very small and very cold, there are huge numbers of grains in a cloud, and each grain emits infrared radiation at long wavelengths.
Some molecules in the
cold gas emit in the
infrared—so, infrared
observations can detect
very cold clouds of gas.

The Interstellar Medium

Also, infrared light penetrates ISM dust better than shorter-wavelength radiation.
So, astronomers can use infrared cameras to look into and through interstellar clouds that are opaque to visible light.

The Interstellar Medium

X-ray observations can detect regions of very hot gas apparently produced by exploding stars.

The Interstellar Medium

Radio observations reveal the emissions of specific molecules in the interstellar medium—the equivalent of emission lines in visible light.
Such studies show that some of the atoms in space have linked together to form molecules.

The Interstellar Medium

There is no shortage of evidence concerning the interstellar medium.
Clearly, space is not empty.
From the correlation between young stars and clouds of gas and dust, you can suspect that stars form from these clouds.

The Interstellar Medium

To study the formation of stars, you can continue comparing theory with evidence.