804 What You’ll Learn • What the structure and processes of the Sun are. • What properties are used to observe and measure stars. • How stars change during their lives and what is left when they die. Why It’s Important The Sun is vital to life on Earth. To understand the Sun, which is a star, it is necessary to understand how all stars function and evolve. Stars are also the building blocks of our galaxy and the source of most elements in the universe. Stars Stars 30 30 NGC 3603 To find out more about the Sun and other stars, visit the Earth Science Web Site at ear thgeu.com
Transcript
804
What You’ll Learn• What the structure and
processes of the Sun are.
• What properties areused to observe andmeasure stars.
• How stars change duringtheir lives and what isleft when they die.
Why It’s ImportantThe Sun is vital to life onEarth. To understand theSun, which is a star, it isnecessary to understandhow all stars function andevolve. Stars are also thebuilding blocks of ourgalaxy and the source of most elements in theuniverse.
StarsStars3030
NGC 3603
To find out more about theSun and other stars, visitthe Earth Science Web Siteat earthgeu.com
Although the Sun is an averagestar, it has many complex processes.We get a glimpse of these processesthrough the solar activity cycle. Theactivity cycle of the Sun refers to howactive its surface is. During the peak ofthe activity cycle, the surface is violentand eruptive and has numerous darkspots. The activity cycle varies with a period of approximately 11 yearsand last peaked in late 2000.
1. Find sites on the Web that providecurrent images of the Sun andobserve how the Sun looks todaywhen viewed at different wave-lengths. Or, go to earthgeu.comand follow the links.
2. Make notes of features that youobserve and the wavelengths fromwhich you observe them.
3. Find out where the Sun cur-rently is in its activity cycle,and compare what you actuallyobserve with what you wouldexpect for this time in thecycle.
Observe Is the Sun near thepeak of its activity cycle? Do thesolar images that you observed fitwith your expectation of wherethe Sun is in its activity cycle?Compare and contrast the imagesobtained at different wavelengths.
Monitoring the SunDiscovery LabDiscovery Lab
OBJECTIVES
• Explore the structure ofthe Sun.
• Describe the solar activ-ity cycle and how the Sunaffects Earth.
• Compare the differenttypes of spectra.
VOCABULARY
photosphere solar flarechromosphere prominencecorona fusionsolar wind fissionsunspot spectrum
Humans have probably always been aware of the Sun. However, it hasbeen only recently that astronomers have begun to understand itsnature. Through observations and probes such as the SolarHeliospheric Observatory (SOHO) and the Ulysses mission, astron-omers have begun to unravel the mysteries of the Sun. Astronomersstill rely on computer models for an explanation of the interior of theSun because the interior cannot be directly observed.
PROPERTIES OF THE SUNThe Sun is the largest object in the solar system, in both size andmass. It would take 109 Earths lined up edge to edge to fit across theSun, or almost 10 Jupiters. The Sun is about 330 000 times as mas-sive as Earth and 1048 times the mass of Jupiter. In fact, the Sun con-tains more than 99 percent of all the mass in the solar system.It should therefore come as no surprise that the Sun’s mass controlsthe motions of the planets and other objects.
The Sun’s average density is similar to the densities of the gas giantplanets, represented by Jupiter in Table 30-1. Astronomers can deducedensities at specific points inside the Sun, as well as other information,only by using computer models that explain the observations thatthey make. These models show that the density in the center of theSun is about 1.50 � 105 kg/m3, which is about thirteen times the den-sity of lead! A pair of dice having this density would weigh about twopounds. However, unlike lead, which is a solid, the solar interior isgaseous throughout because of its high temperature—about 1 � 107 K
in the center. At this high temperature many of thegases are completely ionized, meaning that they arecomposed only of atomic nuclei and electrons. Thisstate of matter is known as plasma. The outer layers ofthe Sun are not quite hot enough to be plasma.
THE SUN’S ATMOSPHEREThe lowest layer of the Sun’s atmosphere, approxi-mately 400 km in thickness, is called the photosphere.This is the visible surface of the Sun, as shown inFigure 30-1A. You may wonder why the photosphereis the visible surface when it is also the lowest layer ofthe Sun’s atmosphere. This is because most of the light
emitted by the Sun comes from thislayer. The two layers above are trans-parent at most wavelengths of visiblelight. Additionally, the top two layersare very dim in the wavelengths thatthey do emit. The average temperatureof the photosphere is about 5800 K.
Above the photosphere is thechromosphere, which is approxi-mately 2500 km in thickness and has atemperature of nearly 30 000 K at thetop. Normally, the chromosphere isvisible only during a solar eclipse,when the photosphere is blocked.
Figure 30-1 The photo-sphere (A) is the visible surface of the Sun. Thisannular eclipse (B) on May30, 1994 shows the redchromosphere of the Sun inthe lower half of the image.
However, astronomers can use special filters to observe the chromo-sphere when the Sun is not eclipsed. The chromosphere appears red,as shown in Figure 30-1B, because it emits most strongly in a narrowband of red wavelengths.
The top layer of the Sun’s atmosphere, called the corona, extendsseveral million kilometers from the top of the chromosphere and hasa temperature range of 1 million to 2 million K. The density of thegas in the corona is very low, which explains why the corona is so dimthat it can be seen only when the photosphere is blocked by eitherspecial instruments, as in a coronagraph, or by the Moon during aneclipse, as in Figure 30-2.
Solar Wind The corona of the Sun does not end abruptly. Instead,gas flows outward from the corona at high speeds and forms the solar wind. As this wind of charged particles, or ions, flows outwardthrough the entire solar system, it bathes each planet in a flood of par-ticles. At 1 AU, Earth’s distance from the Sun, the solar wind flows ata speed of about 400 km/s. The charged particles are deflected byEarth’s magnetic field and are trapped in two huge rings in Earth’smagnetic field called the Van Allen belts. The high-energy particles inthese belts collide with gases in Earth’s atmosphere and cause thegases to give off light that we see as the aurora, shown in Figure 30-3.
Figure 30-2 The faintcorona was a fabulous sightduring the eclipse on July11, 1991 in Baja, California.
Figure 30-3 The aurora isthe result of the particlesfrom the Sun colliding withgases in Earth’s atmosphere.It is most easily viewedfrom regions around thepoles of Earth.
Solar ActivityWhile the solar wind and layers of the Sun’s atmo-sphere are permanent features, other features onthe Sun change over time in a process called solaractivity. The Sun’s magnetic field disturbs the solaratmosphere periodically and causes new features toappear. The most obvious features are sunspots,shown in Figure 30-4, which are dark spots on thesurface of the photosphere. Sunspots are actuallyvery bright, but they appear darker than the sur-rounding areas on the Sun because they are cooler.They are located in regions where the Sun’s intensemagnetic fields poke through the photosphere.These magnetic fields prevent hot gases inside theSun from rising to the surface and heating thespots. Sunspots typically last two months. Theyoccur in pairs with opposite magnetic polarities—with a north and a south pole like a bar magnet.
Solar Activity Cycle Astronomers have observed that the num-ber of sunspots changes regularly and on average, reaches a maximumnumber every 11.2 years on average. Scientists therefore hypothesizedthat the solar activity cycle is 11.2 years in length. However, when thepolarity of the Sun’s magnetic field is taken into account, the lengthof the cycle doubles to 22.4 years. The Sun’s magnetic field reverses,so that the north magnetic pole becomes the south magnetic pole andvice versa. When the polarities of the Sun’s magnetic poles reverse, thepolarities of pairs of sunspots also reverse, because sunspots are
caused by magnetic fields. Thus, the solar activ-ity cycle starts with minimum spots and pro-gresses to maximum spots. Then the magneticfield reverses in polarity, and the spots start at aminimum number and progress to a maximumnumber again. The magnetic field then switchesback to the original polarity and completes thesolar activity cycle.
Other Solar Features Coronal holes, shownin Figure 30-5, are often located over sunspotgroups. Coronal holes are areas of low density inthe gas of the corona. They are the main regionsfrom which the particles that comprise the solarwind escape. These holes are visible in the X rayregion of the electromagnetic spectrum.
808 CHAPTER 30 Stars
Figure 30-4 Sunspots con-sist of two regions: thelighter outer ring called thepenumbra, and the darkinner ring called the umbra.
Figure 30-5 The darkregions in this X-ray imageare coronal holes.
30.1 The Sun 809
Highly active solar flares also are associated with sunspots.Solar flares are violent eruptions of particles and radiation from thesurface of the Sun, as shown in Figure 30-6A. Often, the released par-ticles escape the surface of the Sun in the solar wind and Earth getsbombarded with the particles a few days later. The largest solar flareon record, which occurred in April 2001, hurled particles from theSun’s surface at 7.2 million km/hr.
Another active feature, sometimes associated with flares, is aprominence, shown in Figure 30-6B, which is an arc of gas that isejected from the chromosphere, or gas that condenses in the innercorona and rains back to the surface. Prominences can reach temper-atures greater than 50 000 K and can last from a few hours to a fewmonths. Like flares, prominences also are associated with sunspots,and hence, occurrences of both vary with the solar activity cycle.
Impact on Earth There is evidence that the solar activity cycleaffects climates on Earth. For example, some scientists have found evi-dence of subtle climate variations within 11-year periods. Also, therewere severe weather changes on Earth during the latter half of the1600s when the solar activity cycle stopped and there were no sunspotsfor nearly 60 years. No one knows why the Sun’s cycle stopped. Those60 years were known as the “Little Ice Age” because the weather wasvery cold in Europe and North America during those years.
The Solar InteriorYou may be wondering where all the energy that causes solar activityand light comes from. Within the core of the Sun, where the pressureand temperature are extremely high, fusion occurs. Fusion is thecombining of lightweight nuclei, such as hydrogen, into heaviernuclei. This is the opposite of the process of fission, which is thesplitting of heavy atomic nuclei into smaller, lighter atomic nuclei.
A B
Figure 30-6 This solar flare(A) was observed by theNational Solar Observatory,and this solar prominence(B) was observed by the Big Bear Solar Observatoryof the New Jersey Institute of Technology.
EnvironmentalConnection
In the core of the Sun, helium is a product of the process in whichhydrogen nuclei fuse. The mass of the helium nucleus is less than thecombined mass of the hydrogen nuclei, which means that mass isbeing lost during the process somehow. Albert Einstein’s theory ofspecial relativity showed that mass and energy are equivalent, andthat matter can be converted into energy and vice versa. This rela-tionship can be expressed as E = mc 2, where E is energy measured injoules, m is the quantity of mass that is converted to energy measuredin kilograms, and c is the speed of light measured in m/s. This the-ory explains that the mass lost in the fusion of hydrogen to helium isconverted to energy, which powers the Sun. At the Sun’s rate ofhydrogen fusing, it is about halfway through its lifetime, withanother 5 billion years or so left.
Energy from the Sun The quantity of energy that arrives onEarth every day from the Sun is enormous. Above Earth’s atmo-sphere, 1354 J of energy is received in 1 m2 per second (1354 W/m2).In other words, thirteen 100-W lightbulbs could be operated withthe solar energy that strikes a 1-m2 area. However, not all of thisenergy reaches the ground because some is absorbed and scatteredby the atmosphere. You will learn how energy from the Sun can beconverted to electricity with solar panels in the Problem-Solving Labon this page.
810 CHAPTER 30 Stars
Calculate energy output from solar panels The energy from the Sunstriking the top of Earth’s atmosphere is1354 W/m2, about half of which reachesthe ground. The energy output is mea-sured in watts, because it is energyreceived per second. Assume that solarpanels can convert about 15 percent of the Sun’s energy into electricity.
Analysis1. How much energy from the Sun reaches
the ground per square meter per sec-ond?
2. How much of the Sun’s energy would a1-m2 solar panel receive per second?
3. After the 1-m2 solarpanel converts theenergy to electric-ity, how muchelectricity isavailable persecond?
ThinkingCritically4. How many 1-m2 solar
panels would be needed to produce1000 W of electricity for a house?
5. How would a cloudy region affect theenergy production of solar panels? Isthere anything that could be done tocounteract this? Explain.
Using Numbers
Topic: Farthest StarTo find out more about theSun and other stars, visitthe Earth Science Web Siteat earthgeu.com
Activity: Research themost recent discoveries ofstars. Which star is the mostdistant from Earth? Howmany light years or AUs isthe star from the Earth?
Figure 30-7 It takes about170 000 years to transferenergy to the surface of theSun through these zones.
A
B
Figure 30-8 A continuousspectrum is produced by ahot solid, liquid, or densegas. When a cloud of gas isin front of this hot source,an absorption spectrum isproduced. A cloud of gaswithout a hot sourcebehind it will produce anemission spectrum (A). TheSun’s spectrum (B) is anabsorption spectrum.
Solar Zones If the energy of the Sun is produced in thecore, how does it get to the surface before it travels to Earth?The answer lies in the two zones in the solar interior. Above thecore is a region called the radiative zone, which extendsapproximately 86 percent of the way to the photosphere. Inthis zone, energy is transferred from particle to particle byradiation, as atoms continually absorb energy and then re-emitit. Above the radiative zone is the convective zone. In this zone,moving volumes of gas carry the energy the rest of the way tothe Sun’s surface through convection, which you learned aboutin Chapter 11. The radiative and convective zones are illus-trated in Figure 30-7.
SPECTRAYou are probably familiar with the rainbow that appears when whitelight is shined through a prism. This rainbow is a spectrum, whichis visible light arranged according to wavelengths. There are threetypes of spectra: continuous, emission, and absorption, as shown inFigure 30-8. All three types will be discussed on the next few pages.A spectrum that has no breaks in it, such as the one produced whenlight from an ordinary bulb is shined though a prism, is called a con-tinuous spectrum. A continuous spectrum also can be produced bya glowing solid or liquid, or by a highly compressed, glowing gas.However, if you were to observe a spectrum coming from a non-compressed gas, you would see bright lines at certain wavelengths.This is called an emission spectrum, and the lines are called emissionlines. The wavelengths of the lines you see depend on the elementbeing observed, because each element has its own characteristicemission spectrum.
Radiative zone Convective zone
25% 61% 14%
Conversely, if you observe the light from the Sun ina spectrum, you will see a series of dark bands. Thesedark spectral lines are caused by different chemical ele-ments that absorb light at specific wavelengths. This iscalled an absorption spectrum, and the lines are calledabsorption lines. Absorption is caused by a cooler gasin front of a source that emits a continuous spectrum.The absorption lines caused by the element in the gasare in the exact same location as the emission linesmade by the same element. Thus, by comparing labo-ratory spectra of different gases with the dark lines inthe solar spectrum, it is possible to identify the ele-ments that make up the Sun’s outer layers. You willexperiment with identifying spectral lines in theGeoLab at the end of this chapter.
SOLAR COMPOSITIONThe Sun consists of hydrogen, about 70.4 percent by mass, andhelium, 28 percent, as well as a small amount of other elements, asillustrated in Figure 30-9. This composition is very similar to that ofthe gas giant planets, which suggests that the Sun and the gas giantsrepresent the composition of the interstellar cloud from which thesolar system formed, while the terrestrial planets have lost most ofthe lightweight gases, as you learned in Chapter 29. The Sun’s com-position represents that of the galaxy as a whole. Most stars haveproportions of the elements similar to the Sun. Hydrogen andhelium are the predominant gases in stars, as well as in the entireuniverse. All other elements are in very small proportions comparedto hydrogen and helium.
812 CHAPTER 30 Stars
1. How do astronomers know what condi-tions exist inside the Sun?
2. Describe the layers of gas above the Sun’svisible surface.
3. How does energy produced in the core ofthe Sun reach the surface? How long doesit take?
4. How are the Sun’s magnetic field and itsactivity cycle related?
5. How are the different types of spectracreated?
6. Thinking Critically How would the Sunaffect Earth if Earth did not have a mag-netic field?
SKILL REVIEW
7. Comparing and Contrasting Compare andcontrast solar flares, prominences, andsunspots. For more help, refer to the SkillHandbook.
H 70.4%
He 28%
O 0.756%C 0.278%Ne 0.169%Fe 0.123%N 0.0814%Si 0.0696%Mg 0.0645%S 0.0479%
Composition of the Sun by Mass
Figure 30-9 The Sun is primarily composed ofhydrogen and helium.Trace elements in the Sunthat are not listed are Li,Be, B, F, Al, P, Cl, Ar, K, Ca,Sc, Tr, V, Cr, Mn, Co, Ni, Cu,and Zn.
When you look up at the sky at night, you can often see the bright-est stars, even in the city. These stars appear to be fairly isolated fromeach other. However, away from city lights, you would notice manymore stars grouped together.
GROUPS OF STARSLong ago, many civilizations looked at the brightest stars and namedgroups of them after animals, mythological characters, or everydayobjects. These groups of stars are called constellations. Today, wegroup stars by the 88 constellations named by ancient peoples. Someconstellations can be seen all year long, depending on the observer’slocation. In the northern hemisphere, you can see constellations thatappear to move around the north pole of Earth. These constellationsare called circumpolar constellations. Ursa Major, also known as theBig Dipper, is a circumpolar constellation for the northern hemi-sphere.
Unlike circumpolar constellations, the other constellations can beseen only at certain times of the year because of Earth’s changingposition in its orbit around the Sun, as illustrated in Figure 30-10.For example, the constellation Orion can be seen only in the north-ern hemisphere’s winter, and the constellation Hercules can be seenonly in the northern hemisphere’s summer. This is why constella-tions are classified as summer, fall, winter, and spring constellations.For maps of the constellations, see Appendix K.
Star Clusters Although the stars in constellations appear to beclose to each other, very few are gravitationally bound to one other.The reason that they appear to be close together is that human eyescan’t distinguish how far or near stars actually are. Two stars could
OrionHercules
Northernhemispheresummer
Sun
Northernhemispherewinter
Figure 30-10 Dependingon the time of year, onlycertain constellations arevisible. (not to scale)
814 CHAPTER 30 Stars
appear to be right next to each other, almost touching, but one mightbe 1 trillion km away from Earth, and the other might be 2 trillion kmaway from Earth. However, by measuring distances to stars andobserving how they interact with each other, scientists can determinewhich stars are gravitationally bound to each other. A group of starsthat are gravitationally bound to each other is called a cluster. ThePleiades, in the constellation Taurus, shown in Figure 30-11A, is anopen cluster because the stars are not densely packed. In contrast, aglobular cluster is a group of stars that are densely packed into aspherical shape, such as M13, in the constellation Hercules, shown inFigure 30-11B.
Binaries When only two stars are gravitationally bound togetherand orbit a common center of mass, they are called a binary star.More than half of the stars in the sky are either binary stars or mem-bers of multiple-star systems. The bright star Sirius is actually abinary system, as shown in Figure 30-12. Most binary stars appearto be single stars to the human eye, even with a telescope. The twostars are usually too close together to appear separately and one ofthe two is often much brighter than the other. Astronomers are ableto identify binary stars through several methods. For example, evenif only one star is visible, accurate measurements can show that itsposition shifts back and forth as it orbits the center of mass betweenit and the unseen companion star. Also, the orbital plane of a binarysystem can sometimes be seen edge-on from Earth. In such cases,the two stars alternately block each other out and cause the totalbrightness of the two-star system to dip each time one star eclipsesthe other. This type of binary star is called an eclipsing binary.
Figure 30-11 The Pleiades(A) is an open cluster, andM13 (B) is a densely packedglobular cluster.
A B
Figure 30-12 The brightstar Sirius and its compan-ion white dwarf, on theleft, are a binary system.
30.2 Measuring the Stars 815
STELLAR POSITIONS AND DISTANCESAstronomers use two units of measure for long distances. One,which you are probably familiar with, is a light-year (ly). A light-yearis the distance that light travels in one year, equal to 9.461 � 1012 km.Astronomers often use a larger unit than a light-year, a parsec. A par-sec (pc) is equal to 3.26 ly, or 3.086 � 1013 km.
Precise position measurements are an important tool for findingdistances to stars. To estimate the distance of stars from Earth,astronomers make use of the fact that nearby stars shift in positionas observed from Earth. This apparent shift in position caused by themotion of the observer is called parallax. In this case, the motion ofthe observer is the change in position of Earth as it orbits the Sun. AsEarth moves from one side of its orbit to the opposite side, a nearbystar appears to be shifting back and forth, as illustrated in Figure 30-13.The closer the star, the larger the shift. The distance to a star can beestimated from its parallax shift. In fact, a parsec is defined as the dis-tance at which an object has a parallax of 1 arcsecond. Using the par-allax technique, astronomers only could find accurate distances tostars up to about 100 pc, or approximately 300 ly, away until recently.Now with advancements in technology, such as the Hipparcos satel-lite, astronomers can find accurate distances up to 500 pc by usingparallax. In the MiniLab later in this section, you will experimentwith distance and how it affects parallax shifts.
Basic Properties of StarsThe basic properties of stars include diameter, mass, brightness,energy output (power), surface temperature, and composition. Thediameters of stars range from as little as 0.1 times the Sun’s diameterto hundreds of times larger, while their masses vary from a little lessthan 0.01 to 20 or more times the Sun’s mass. The most massive starscan be as massive as 50 to 100 Suns but are extremely rare.
July
Background stars
Nearbystar
Earth in July
Sun
Earth in January
January
Figure 30-13 The shift inposition of a star as viewedfrom opposite sides ofEarth’s orbit around theSun is called parallax.
Magnitude One of the most basic observable properties of a staris how bright it appears. The ancient Greeks established a classifica-tion system based on the brightnesses of stars. The brightest starswere given a ranking of +1, the next brightest +2, and so on. Today’sastronomers still use this system, but they have refined it.
Apparent Magnitude Astronomers have defined the ancientGreek system of classification as apparent magnitude, or how brighta star appears to be. In this system, a difference of 5 magnitudes corre-sponds to a factor of 100 in brightness. Thus, a magnitude +1 star is100 times brighter than a magnitude +6 star. A difference of 1 magni-tude corresponds to a factor of 2.512 in brightness. The modern mag-nitude system extends to objects that are both brighter and fainterthan those that were included in the ancient Greek system. For objectsbrighter than magnitude +1, such as the Sun, the Moon, Venus, andsome of the very brightest stars, negative numbers are assigned. Theapparent magnitudes of several objects are shown in Figure 30-14.
Absolute Magnitude Apparent magnitude does not actuallyindicate how bright a star is, because it does not take distance intoaccount. A faint star can appear to be very bright because it is rela-tively close to Earth, while a bright star can appear to be faint becauseit is far away. To account for this phenomenon, astronomers havedeveloped another classification system for brightness. Absolutemagnitude is the brightness an object would have if it were placed ata distance of 10 pc. The classification of stars by absolute magnitudeallows comparisons that are based on how bright the stars wouldappear at equal distances from an observer. The absolute magnitudesfor several objects are shown in Figure 30-14. The disadvantage ofabsolute magnitude is that it can be calculated only when the actualdistance to a star is known.
Figure 30-14 This showsthe apparent and absolutemagnitudes of some famil-iar celestial objects.
Using Numbers Thedifference in bright-ness between a mag-nitude +12 star and amagnitude +9 star is2.512 � 2.512 �2.512 = 2.5123 =15.85. What is thedifference in bright-ness between a mag-nitude +21 star and amagnitude +14 star?
Luminosity Apparent magnitudes do notgive an actual measure of energy output. Tomeasure the energy output from the surface ofa star per second, called its luminosity, anastronomer must know both the star’s appar-ent magnitude and how far away it is. Thebrightness we observe for a star depends onboth its luminosity and its distance, andbecause brightness diminishes with the squareof the distance, a correction must be made fordistance. Luminosity is measured in units ofenergy emitted per second, or watts. The Sun’sluminosity is about 3.85 � 1026 W. This isequivalent to 3.85 � 1024 100-W lightbulbs!The values for other stars vary widely, fromabout 0.0001 to more than a million times theSun’s luminosity. No other stellar propertyvaries by so much.
Spectra of StarsYou have learned that the Sun has darkabsorption lines at specific wavelengths inits spectrum. Other stars also have darkabsorption lines in their spectra and areclassified according to their patterns ofabsorption lines.
Classification Stars are assigned spectraltypes in the following order: O, B, A, F, G, K,and M. Each class is subdivided into morespecific divisions with numbers from 0 to 9.For example, a star may be classified as beinga type A4 or A5. The classes were originallybased only on the pattern of spectral lines,but astronomers later discovered that theclasses correspond to stellar temperatures,with the O stars being the hottest and the Mstars being the coolest. Thus, by examinationof a star’s spectrum, it is possible to estimateits temperature. The Sun is a type G2 star,which corresponds to a surface temperatureof about 5800 K. Surface temperatures rangefrom about 50 000 K for the hottest O stars toas low as 2000 K for the coolest M stars.
30.2 Measuring the Stars 817
Parallax in theClassroom
Model stellar parallaxand the change in parallax angle with distance.
Procedure1. Place a meterstick at a fixed position and
attach a 4-m piece of string to each end.2. Stand away from the meterstick and hold
the two strings together to form a trian-gle. Be sure to hold the strings taut.Measure your distance from the meter-stick. Record your measurement.
3. Measure the angle between the twopieces of string with a protractor. Recordyour measurement of the angle.
4. Repeat steps 2 and 3 for different dis-tances from the meterstick by shorteningor lengthening the string.
5. Make a graph of the angles versus theirdistance from the meterstick.
Analyze and Conclude1. What does the length of the meterstick
represent? The angle?2. What does the graph show? How does
parallax angle depend on distance?3. Are the angles that you measured similar
to actual stellar parallax angles? Explain.
818 CHAPTER 30 Stars
Figure 30-15 These are the typical black and whiteabsorption spectra of a classB5 star (A), a class F5 star(B), a class K5 star (C), anda class M5 star (D).
A
B
C
D
All stars, including the Sun, have nearly identical compositions,despite the differences in their spectra, shown in Figure 30-15. Thedifferences in the appearance of their spectra are almost entirely aresult of temperature effects. Hotter stars have fairly simple spectrawhile cooler stars have spectra with more lines. The coolest stars havebands in their spectra due to molecules, such as titanium oxide, intheir atmospheres. Typically, about 73 percent of a star’s mass ishydrogen, about 25 percent is helium, and the remaining 2 percent iscomposed of all the other elements. While there are some variationsin the composition of stars, particularly in that final 2 percent, allstars have this general composition.
Wavelength Shifts Spectral lines provide other informationabout stars in addition to composition and temperature. Spectrallines are shifted in wavelength by motion between the source of lightand the observer. The shifts in spectral lines are an example of theDoppler effect.
Unshifted light from star Unshifted light from star
Blueshiftedlight from star
Redshiftedlight from star
1
1234
Motion of star
Figure 30-16 As a starmoves towards or awayfrom an observer, light isblueshifted or redshifted,respectively.
30.2 Measuring the Stars 819
If a star is moving toward the observer, the spectral lines are shiftedtoward shorter wavelengths, or blueshifted. However, if the star is mov-ing away, the wavelengths become longer, or redshifted, as illustrated inFigure 30-16. The higher the speed, the larger the shift, and thus care-ful measurements of spectral line wavelengths can be used to determinethe speed of a star’s motion.
Because there is no Doppler shift for motion that is sideways tothe line of sight, astronomers can learn only about the portion of astar’s motion that is directed toward or away from Earth. TheDoppler shift in spectral lines can be used to detect binary stars asthey move about their center of mass towards and away from Earthwith each orbit.
H-R Diagrams The properties ofmass, luminosity, temperature, anddiameter are closely related. Eachclass of star has a specific mass,luminosity, magnitude, tempera-ture, and diameter. These relation-ships can be demonstrated on agraph called the Hertzsprung-Russell diagram (H-R diagram)with absolute magnitude plottedon the vertical axis and tempera-ture or spectral type plotted on thehorizontal axis, as shown inFigure 30-17. This graph was firstplotted in the early twentieth cen-tury. An H-R diagram with lumi-nosity plotted on the vertical axislooks very similar to the one inFigure 30-17.
About 90 percent of stars,including the Sun, fall along abroad strip of the H-R diagramcalled the main sequence, whichruns diagonally from the upper-leftcorner, where hot, luminous starsare represented, to the lower-rightcorner, where cool, dim stars arerepresented. The interrelatedness ofthe properties of these stars indi-cates that all these stars have similarinternal structures and functions.
Whitedwarfs
Giants
Supergiants
Main sequence
O5+15
+10
+5
Sun
Ab
solu
te m
agn
itu
de
0
–5
B0 B5 A0 A5 F0 F5
Spectral type
G0 G5 K0 K5 M0 M5
40 000 10 000 7 000 6 000 5 000 3 000
Surface temperature (K)
Figure 30-17 This H-R dia-gram shows the relationshipbetween absolute magni-tude, surface temperature,and spectral type. Massdecreases from left to right,while luminosity increasesfrom bottom to top.
820 CHAPTER 30 Stars
Table 30-2 summarizes the basic properties of main sequencestars. But what about the stars that do not lie on the main sequence?The stars plotted at the upper right of the H-R diagram are cool, yetvery luminous. Because cool surfaces emit much less radiation persquare meter than hot ones do, these cool stars must have large sur-face areas to be so bright. For this reason, these larger, cool, luminousstars are called red giants. Red giants are so large—more than 100times as large as the Sun in some cases—that Earth would be swal-lowed up if the Sun were to become a red giant! Conversely, the dim,hot stars plotted in the lower-lefthand corner of the H-R diagrammust be very small, or else they would be far more luminous. Thesesmall, dim, hot stars are called white dwarfs. A white dwarf is aboutthe size of Earth but has a mass about as large as the Sun’s. You willlearn how all the different stars are formed in the following section.
1. Describe two types of stars that are noton the main sequence.
2. Explain the difference between apparentand absolute magnitudes.
3. What are the main properties of stars?
4. How do astronomers know that somestars are binary stars?
5. Thinking Critically How do the Sun’s prop-erties compare with those of other stars?
SKILL REVIEW
6. Interpreting Diagrams Use the H-R dia-gram in Figure 30-17 to describe theproperties of an A and an M star on themain sequence. For more help, refer tothe Skill Handbook.
Table 30-2 Properties of Main-Sequence Stars
Spectral Surface Type Mass* Temperature (K) Luminosity* Radius*
O5 40.0 40 000 5 � 105 18.0
B5 6.5 15 500 800 3.8
A5 2.1 8500 20 1.7
F5 1.3 6580 2.5 1.2
G5 0.9 5520 0.8 0.9
K5 0.7 4130 0.2 0.7
M5 0.2 2800 0.008 0.3
*These properties are given relative to the Sun. For example, an O5 star has a mass 40 times that of the Sun,or 40 � 1.99 � 1030 kg = 7.96 � 1031 kg. The mass, luminosity, and radius of the Sun are 1.99 � 1030 kg,3.846 � 1026 W, and 6.96 � 105 km, respectively.
• Explain how astronom-ers learn about the inter-nal structure of stars.
• Describe how the Sunwill change during itslifetime and how it willend up.
• Compare the evolutionsof stars of differentmasses.
VOCABULARY
nebulaprotostarneutron starsupernovablack hole
30.330.3 Stellar Evolution
Using observations as their guide, astronomers have developed mod-els of stars that successfully explain the properties that you have justlearned about. These models, like the solar model, are based on equa-tions describing physical processes that occur inside stars. However,these models are accepted only when they reproduce the externalproperties that have been observed.
BASIC STRUCTURE OF STARSMass governs a star’s temperature, luminosity, and diameter. In fact,astronomers have discovered that the mass and the composition ofa star determine nearly all its other properties. The more massive astar is, the greater the gravity pressing inward, and the hotter anddenser the star must be inside to balance gravity. The temperatureinside a star governs the rate of nuclear reactions, which in turndetermines the star’s energy output, or luminosity. The balancebetween gravity squeezing inward and pressure from nuclear fusionand radiation pushing outward, called hydrostatic equilibrium,must hold for any stable star; otherwise, the star would expand orcontract, as illustrated in Figure 30-18. This balance is governed bythe mass of a star.
Fusion Inside a star, conditions vary in much the same way thatthey do inside the Sun. The density and temperature increase towardthe center, where energy is generated by nuclear fusion. Stars on themain sequence all produce energy by fusing hydrogen into helium,as the Sun does. Stars that are not on the main sequence either fusedifferent elements in their cores or do not undergofusion at all.
Fusion reactions involving elements other thanhydrogen can occur. Once a star’s core has been con-verted into helium, the helium may fuse to form carbonif the temperature is high enough. At even higher tem-peratures, carbon can react with helium to form oxygen,then neon, then magnesium, and then silicon. Othertypes of reactions can produce even heavier elements,but few heavier than iron. Each of these reactions pro-duces energy according to the equation E = mc 2 as asmall fraction of mass is converted into energy. Thisenergy stabilizes a star by producing the pressure neededto counteract gravity.
Pressure fromfusion andradiation
Gravity
Figure 30-18 This star isstable and will not expandor contract.
822 CHAPTER 30 Stars
Infalling material
Rotating disk
Protostar
A
To learn more about theHubble Space Telescope,go to the NationalGeographic Expeditionon page 902.
STELLAR EVOLUTION AND LIFE CYCLESA star changes as it ages because its internal compositionchanges as nuclear fusion reactions in the star’s core convert one element into another. As a star’s core compo-sition changes, its density increases, its temperature rises,and its luminosity increases. Eventually, the nuclear fuelruns out. Then the star’s internal structure and mechanismfor producing pressure must change to counteract gravity.
Star Formation All stars form in much the same manner as theSun did. The formation of a star begins with a cloud of interstellar gasand dust, called a nebula (pl. nebulae), which collapses on itself as aresult of its own gravity. As the cloud contracts, its rotation forces itinto a disk shape with a hot condensed object at the center, called aprotostar, as illustrated in Figure 30-19A. The condensed object willbecome a new star. A protostar is brightest at infrared wavelengths.
Fusion Begins Eventually, the temperature inside a protostarbecomes hot enough for nuclear fusion reactions to begin. The firstreaction to ignite is always the conversion of hydrogen to helium.Once this reaction begins, the star becomes stable because it then hassufficient internal heat to produce the pressure needed to balancegravity. The object is then truly a star and takes its place on the mainsequence according to its mass. A new star often illuminates the gasand dust surrounding it, as shown in Figure 30-19B.
THE SUN’S LIFE CYCLEWhat happens during a star’s life cycle depends on its mass. Forexample, as a star like the Sun converts hydrogen into helium in itscore, it gradually becomes more luminous because the core densityand temperature rise slowly and increase the reaction rate. It takes
B
Figure 30-19 A protostar,formed from a disk of gasand dust (A), will become a star when fusion begins.The Triffid Nebula (B) isilluminated by new stars, as shown by the HubbleSpace Telescope.
about 10 billion years for a star with the mass of the Sun toconvert all of the hydrogen in its core into helium. Thus, sucha star has a main sequence lifetime of 10 billion years.
Only the innermost 10 percent or so of a star’s mass canundergo reactions because temperatures outside of this corenever get hot enough for reactions to occur. Thus, when thehydrogen in its core is gone, a star has a helium center andouter layers made of hydrogen-dominated gas. Some hydro-gen continues to react in a thin layer at the outer edge of thehelium core, as illustrated in Figure 30-20. The energy pro-duced in this layer forces the outer layers of the star to expandand cool. The star then becomes a red giant, because its lumi-nosity increases while its surface temperature decreases due tothe expansion.
While the star is a red giant, it loses gas from its outer layers. The staris so large that its surface gravity is very low and thus the outer layers canbe driven away by small expansions and contractions or pulsations, ofthe star due to instability. Meanwhile, the core of the star becomes hotenough, at 100 million K, for helium to react and form carbon. The starcontracts back to a more normal size, where it again becomes stable forawhile. The helium-reaction phase lasts only about one-tenth as longas the earlier hydrogen-burning phase. Afterwards, when the helium inthe core is all used up, the star is left with a core made of carbon.
A Nebula Once Again A star of the Sun’s mass never becomeshot enough for carbon to react, so the star’s energy production endsat this point. The outer layers expand once again and are driven offentirely by pulsations that develop in the outer layers. This shell of gasis called a planetary nebula. It has nothing to do with planets, despiteits name. In the center of a planetary nebula, shown in Figure 30-21,the core of the star becomes exposed as a small, hot object about thesize of Earth. The star is then a white dwarf made of carbon.
Pressure in White Dwarfs A white dwarf is stable despite thelack of nuclear reactions because it is supported by the resistance ofelectrons being squeezed close together, and does not require asource of heat to be maintained. This pressure counteracts gravityand can support the core as long as the mass of the remaining core isless than about 1.4 times the mass of the Sun.
A star that has less mass than that of the Sun has a similar life cycle,except that helium may never form carbon in the core, and the starends as a white dwarf made of helium. The main sequence lifetime ofsuch a star is much longer, however, because low-mass stars are dimand do not use up their nuclear fuel very rapidly.
Helium core
Hydrogenfusing in a shell
Figure 30-20 After hydro-gen fusing is done in thecore, it continues in a shell.(not to scale)
Figure 30-21 NGC 6751 isa planetary nebula. Thisimage was taken by theHubble Space Telescope.The white dwarf is thewhite object at the centerof the nebula.
LIFE CYCLES OF MASSIVE STARSFor stars more massive than the Sun, evolution is verydifferent. A more-massive star begins its life in thesame way, but much higher on the main sequence,with hydrogen being converted to helium. However,the star’s lifetime in this phase is short, because thestar is very luminous and uses up its fuel quickly.
A massive star undergoes many more reactionphases and thus produces a rich stew of many ele-ments in its interior. The star becomes a red giant sev-eral times as it expands following the end of eachreaction stage. As more shells are formed by the fusionof different elements, illustrated in Figure 30-22, thestar expands to a larger size and becomes a supergiant,such as Betelgeuse in the Orion constellation.
A massive star loses much of its mass during its lifetime as gasdrifts from its outer layers, or is driven away by a stellar wind. A starthat begins with as many as 8 times the Sun’s mass may end up as awhite dwarf with a final mass less than 1.4 times the Sun’s mass. Thecomposition of a white dwarf is determined by how many reactionphases the star went through before it stopped reacting altogether.Thus, there can be white dwarfs made of oxygen, white dwarfs madeof neon, and so on.
Supernovae Some stars do not lose enough mass to becomewhite dwarfs. A star that begins with a mass between about 8 and 20times the Sun’s mass will end up with a core that is too massive tobe supported by electron pressure. Such a star comes to a very vio-lent end. Once reactions in the core of the star have created iron, nofurther energy-producing reactions can occur, and the core of thestar violently collapses in on itself, as illustrated in Figure 30-23A.As it does so, protons and electrons in the core merge to form
824 CHAPTER 30 Stars
Fe/Ni
Si ➛ Fe, NiO ➛ Si, S
Ne ➛ O, MgC ➛ Ne, Mg
He ➛ C, O
H ➛ He
Core
Figure 30-22 A massive starcan have many shells fusingdifferent elements. Massivestars are responsible for pro-ducing heavier elements.(not to scale)
Infalling material
Core(whitedwarf)
Infallingmaterialrebounds
Core(neutron
star)
Shockwaves
Material explodes outward
Core
A B C
Figure 30-23 The core ofthe star collapses (A) form-ing a neutron star (B), onwhich infalling materialbounces off and causes asupernova (C). (not to scale)
neutrons. Like electrons, neutrons can’t be squeezed too closelytogether. Their resistance to being squeezed creates a pressurethat halts the collapse of the core, and the core becomes a neutron star, as illustrated in Figure 30-23B. A neutron starhas a mass of 1.5 to 3 times the Sun’s mass but a radius of onlyabout 10 km! The density is incredibly high—about 100 tril-lion times more dense than water—and is comparable to thatof an atomic nucleus.
A neutron star forms quickly while the outer layers of the starare still falling inward. This infalling gas rebounds when it strikesthe hard surface of the neutron star as illustrated in Figure 30-23B,and explodes outward, as illustrated in Figure 30-23C. The entireouter portion of the star is blown off in a massive explosion called asupernova (pl. supernovae), shown in Figure 30-24. This explosioncreates elements that are heavier than iron and enriches the universe.
Black Holes Some stars are too massive even to form neutronstars. The pressure from the resistance of neutrons being squeezedtogether cannot support the core of a star if the star’s mass is greaterthan about 3 times the mass of the Sun. A star that begins with morethan about 20 times the Sun’s mass will end up above this mass limit,and it cannot form a neutron star. The resistance of neutrons tobeing squeezed is not great enough to stop the collapse and the coreof the star simply continues to collapse forever, compacting matterinto a smaller and smaller volume. The small, but extremely dense,object that remains is called a black hole because its gravity is soimmense that nothing, not even light, can escape it. You will learnabout the search for black holes in the Science in the News feature atthe end of this chapter.
30.3 Stellar Evolution 825
B
A
Figure 30-24 The topimage (A), shows theregion before Supernova1987A in the LargeMagellanic Cloud, while the bottom image (B)shows the supernova in full bloom.
1. How do astronomers learn about theinternal structure and evolution of stars?
2. How does a new star form?
3. What causes a supernova to occur?Explain.
4. Is the lifetime of a massive star shorter orlonger than a star like the Sun? Why?
5. In what ways is the evolution of a massivestar similar to the evolution of the Sun,and in what ways is it different?
6. Thinking Critically How would the uni-verse be different if massive stars did notexplode at the ends of their lives?
SKILL REVIEW
7. Comparing and Contrasting Compare andcontrast how pressure and gravity are bal-anced or not balanced in main-sequencestars, white dwarfs, neutron stars, andblack holes. For more help, refer to theSkill Handbook.
A n astronomer studying a star or other type of celestialobject often starts by identifying the lines in the object’s
spectrum. The identity of the spectral lines gives astronomersinformation about the chemical composition of the distantobject, along with data on its temperature and other properties.
ProblemIdentify stellar spectral lines based ontwo previously identified lines.
Materialsruler
ObjectivesIn this Geolab, you will:• Develop a scale based on the separa-
tion between two previously identifiedspectral lines.
• Measure wavelengths of spectral lines.• Compare measured wavelengths to
known wavelengths of elements todetermine composition.
Preparation
1. Measure the distance between the twoidentified spectral lines on star 1. Besure to use units that are small enoughto get accurate measurements.
2. Calculate the difference in wave-lengths between the two identifiedspectral lines.
3. Set up your scale by dividing the dif-ference in wavelengths by the mea-sured distance between the twoidentified spectral lines. This willallow you to measure wavelengthsbased on your distance measurementunit. For example, 1 mm = 12 nm.
4. Measure the distance to spectral linesfrom one of the two previously iden-tified spectral lines.
5. Convert your distances to wave-lengths using your scale. You havemeasured the difference in wave-length. This difference must be addedor subtracted to the wavelength ofthe line you measured from. If theline you measured from is to theright of the line you are identifying,then you must subtract. Otherwise,you add.
6. Compare your wavelength measure-ments to the table of wavelengthsemitted by elements, and identify theelements in the spectrum.
7. Repeat this procedure for star 2.
Procedure
GeoLab 827
1. What elements are present in the stars?2. How does your list of elements com-
pare with the list of elements seen inthe periodic table in Appendix G?
3. Can you see any clues in the star’sspectrum about which elements aremost common in the stars? Explain.
Analyze
Conclude & Apply
1. Do both stars contain the same linesfor all the elements in the table?
2. You should notice that some absorp-tion lines are wider than others. Whatare some possible explanations for this?
3. How do the thicker absorption linesof some elements in a star’s spectrumeffect the accuracy of your measure-ments? Is there a way to improve yourmeasurements? Explain.
4. Using the following formula, calculate the percent deviation for 5 of your measured lines.
difference from Percent =
accepted value � 100deviation accepted value
Is there a value that has a high percent deviation? If so, what are some possible explanations for this?
Chandra was originally called the Advanced X-ray Astrophysics Facility. Go toearthgeu.com to find links to more informa-tion on how the observatory received itsnew name. Who was it named for? Why?
Activity
Chandra is in a unique position to help scientistssolve some of astronomy’s most baffling puzzles.
Black HolesChandra has turned its eye toward black
holes—high-gravity objects that can suck inentire stars. Black holes are thought to exist atthe center of galaxies. They may be the source ofthe astounding amounts of energy and radiationemitted by galactic centers. Because nothing, noteven light, can escape from black holes, theycannot be observed directly. However, when mat-ter is pulled into a black hole, it is heated toincredibly high temperatures. The matterbecomes so hot that it emits X rays. Chandra isstudying these X-ray emissions, documenting par-ticles of matter up until the very millisecondbefore they disappear into the black hole. Usingthis information, scientists hope to learn moreabout the nature of black holes and the energythey produce. The information Chandra gatherspromises to change our views of the universe.
828 CHAPTER 30 Stars
Chandra is the third of NASA’s Great Observatories—space telescopes designed to capture images beyond the reach of Earth-basedtelescopes. Unlike the Hubble Space Telescopeand the Compton Gamma-Ray Observatory, whichtake pictures created by visible light and gammarays, Chandra studies X rays. These rays areabsorbed by Earth’s atmosphere. Thus, they are best studied from a position high above our planet.
The Keen EyeX-ray telescopes are not new. However,
Chandra represents a vast improvement in tech-nology. According to NASA, Chandra is a billiontimes more powerful than the first X-ray telescope,built just decades ago. “Chandra’s resolvingpower,” a NASA document explains, “is equivalentto the ability to read a 1- cm newspaper headlineat the distance of a half-mile.”
The $2.8 billion observatory is composed ofthree separate elements. The spacecraft systemcontains computers, data recorders, and commu-nication equipment to transmit information backto Earth. The telescope system includes anassembly of high-resolution mirrors, the largestand most polished of their kind. The scienceinstrument system is equipped with a high-resolution camera to record X-ray images of tur-bulent, high-temperature events, such as super-novae. Armed with this space-age technology,
Chandra—An Eye on the UniverseSome of the hottest action in the universe—action that takes place in
black holes, exploding stars, and colliding galaxies—can’t be observed
with the naked eye. With the 1999 launch of the Chandra X-ray
Observatory, however, scientists now have an unprecedented view.Supernova remnant
Main Ideas• The Sun contains most of the mass in the solar system and is
made up primarily of hydrogen and helium.• Astronomers learn about conditions inside the Sun by a combi-
nation of observation and theoretical models.• The Sun’s atmosphere consists of the photosphere, the chromo-
sphere, and the corona.• The Sun has a 22-year activity cycle caused by reversals in its
magnetic field polarities.• Sunspots, solar flares, and prominences are active features of the
Sun.• The solar interior consists of the core, where fusion of hydrogen
into helium occurs, and the radiative and convective zones.
Main Ideas• Positional measurements of the stars are important for measur-
ing distances through stellar parallax shifts.• Stellar brightnesses are expressed in the systems of apparent
and absolute magnitude.• Stars are classified according to the appearance of their spectra,
which indicate the surface temperatures of stars.• The H-R diagram relates the basic properties of stars: class,
mass, temperature, and luminosity.
Main Ideas• The mass of a star determines its internal structure and its other
properties.• Gravity and pressure balance each other in a star.• If the temperature in the core of a star becomes high enough,
elements heavier than hydrogen but lighter than iron can fusetogether.
• Stars such as the Sun end up as white dwarfs. Stars up to about8 times the Sun’s mass also form white dwarfs after losing mass.Stars with masses between 8 and 20 times the Sun’s mass endas neutron stars, and more massive stars end as black holes.
• A supernova occurs when the outer layers of the star bounce offthe neutron star core, and explode outward.
SECTION 30.1
The Sun
SECTION 30.2
Measuringthe Stars
SECTION 30.3
Stellar Evolution
Study Guide 829earthgeu.com/vocabulary_puzzlemaker
1. Which of the following is not a part of the Sun’satmosphere?a. the corona c. the solar windb. the chromosphere d. the photosphere
2. Which of the following is not created by the Sun’smagnetic field?a. the radiative zone c. solar flaresb. prominences d. sunspots
3. Which type of spectrum, if any, does the Sun emit?a. an absorption spectrumb. an emission spectrumc. a continuous spectrumd. no spectrum
4. Where does nuclear fusion in the Sun occur?a. in the convective zoneb. in the radiative zonec. in the photosphered. in the core
5. If a star begins its evolution with 10 times themass of the Sun, but ends with 2 times the massof the Sun, what type of object does it form?a. a white dwarf c. a neutron starb. a nebula d. a black hole
6. How would you calculate how much brighter amagnitude +4 star is than a magnitude +7 star?a. 2.512 � 2.512b. 2.512 � 2.512 � 2.512c. 2.512 � 2.512 � 2.512 � 2.512d. 2.512 � 2.512 � 2.512 � 2.512 � 2.512
7. Which of the following is the correct order ofunits from largest to smallest?a. pc, ly, AU, km c. km, pc, ly, AUb. ly, pc, km, AU d. km, AU, pc, ly
Understanding Main Ideas 8. What is the difference between nuclear fusionreactions and nuclear fission reactions?
9. Why do we say that the Sun’s activity cycle is 22 years long, when the number of sunspots follows an 11-year pattern?
10. Why do sunspots appear to be darker than theirsurroundings?
11. What is the energy source for all main-sequencestars?
12. Does the Doppler shift provide complete informa-tion about the motion of a star? Explain.
Use the diagram below to answer question 13.
13. Use the following terms to fill in the concept mapof the evolution of a star like the Sun.
star nebula red giantprotostar white dwarf planetary nebula
SLOW DOWN Read the questions and answerchoices carefully. Remember that doing most ofthe problems and getting them right is alwayspreferable to doing all the problems and gettingmany of them wrong.
