Download Chapter 11 - Astronomy

CHAPTER 11
The Sun
CHAPTER OUTLINE
11-1 Solar Properties
1. As viewed from the Earth, the Sun has an average angular diameter of 31´59° and is at
an average distance from the Earth of 1.50  108 km. The diameter of the Sun can therefore be calculated as 1.39  106 km, or about 110 times Earth’s diameter and about 10
times Jupiter’s.
2. Using Kepler’s 3rd law we find that the mass of the Sun is 1.99  1030 kg, about
333,000 times Earth’s mass. Thus, the Sun’s average density is 1.41 g/cm3, about the
same as the density of Jupiter.
3. By observing sunspots, Galileo concluded that the Sun’s rotation period is more than a
month. Today we know that the Sun exhibits differential rotation (25 days at its equator
and nearly 35 days near its poles).
4. Sunspots are regions of the photosphere that are temporarily cool and dark compared
to surrounding regions.
11-2 Solar Energy
1. The Sun emits energy in all parts of the electromagnetic spectrum.
2. Luminosity is the rate at which electromagnetic energy is emitted; the solar luminosity
is about 3.85  1026 watts.
3. Solar energy strikes the Earth at the rate of 1,370 watts/m2.
The Source of the Sun’s Energy
1. In the mid-1800s, it was proposed by von Helmholtz and Kelvin that a slow gravitational contraction is the source of the Sun’s energy. Such a mechanism would power the
Sun for only a few hundred million years.
2. By the early 20th century, however, geologists showed that the Earth was billions of
years old—a period 10 times longer than that necessary for gravitational contraction to
produce the Sun’s energy. The theory had to be abandoned.
3. Einstein proposed in 1905 that mass and energy are interconvertible. During the 1930s,
Hans Bethe worked out the theory of solar nuclear reactions.
Solar Nuclear Reactions
1. An atom’s nucleus is composed of protons (massive, positively charged particles) and
neutrons (massive nuclear particles with no electric charge).
2. Nuclear reactions involve forces between nuclear particles.
3. In nuclear fusion, two nuclei combine to form a larger nucleus, releasing energy in the
process. In the Sun, 4 hydrogen nuclei are fused to form 1 helium nucleus and energy.
4. The fusion process in the Sun converts only 0.7% of the original mass into energy.
5. To produce the Sun’s energy output, nearly 5 million metric tons of matter must be
converted into energy each second. This, in turn, requires 626 billion kg of hydrogen be
transformed into 622 billion kg of helium.
6. The proton-proton chain is the series of nuclear reactions that begins with 4 protons
and ends with a helium nucleus. This chain is the main fusion process in the Sun, responsible for 98.5% of the energy production.
7. Deuterium is a hydrogen nucleus that contains one neutron and one proton.
8. A positron is a positively charged electron emitted from the nucleus in some nuclear
reactions.
9. A neutrino is an elementary particle that has little mass and no charge but carries energy from a nuclear reaction.
10. Fusion can only occur if the nuclei are moving fast enough (i.e. at high temperature)
to overcome the electrical repulsion that makes them repel each other.
Advancing the Model: Fission and Fusion Power on Earth
1. The nuclear reaction in the H-bomb is similar to that in the Sun; it produces helium
from hydrogen.
2. Peaceful uses of fusion power have not yet been developed, although much research
has taken place over the last 6 decades and is still in progress.
3. Fission involves the release of energy when a large nucleus is broken into two medium
nuclei, with mass being converted into energy in the process.
4. The atomic bomb was the first application of fission power. Today we use the energy
of fission to produce electricity in nuclear power plants.
11-3 The Sun’s Interior
1. Because of the high temperatures on and in the Sun, most electrons are stripped away
from their nuclei. Most of the material of the Sun’s interior consists of free nuclei and
free electrons.
2. The behavior of this solar material is similar to that of a simple gas. The most important properties in describing gases are: temperature, pressure, and particle density.
3. Particle density is the number of separate atomic and/or nuclear particles per unit of
volume.
Pressure, Temperature, and Density
1. Pressure is the force per unit area. The two factors that determine the pressure of a gas
are its density and the speed of its constituent particles.
2. Since the temperature of a gas is a measure of the speed of its constituent particles,
pressure is related to temperature.
3. The greater the density of a gas in a container, the higher the number of collisions with
the walls of the container. Thus, pressure is related to density.
4. Under ideal conditions, gas pressure is proportional to its density times its temperature.
Hydrostatic Equilibrium
1. In a star or a planet, hydrostatic equilibrium is the balance between pressure caused
by the weight of material above a thin layer and the upward pressure exerted by material
below that layer.
2. Pressure at the Sun’s center is calculated to be 2.5  1011 times that on the surface of
the Earth.
3. The high pressure and temperature at the Sun’s core results in fusion reactions; the solar core extends out to about 25% of the Sun’s radius.
Energy Transport
1. Three possible methods by which energy can be transferred from the center of the Sun
outward:
(a) Conduction: The transfer of energy in a solid by collisions between atoms and/or
molecules. This is not a significant factor in transporting energy within the Sun.
(b) Convection: the transfer of energy in a gas or liquid by means of the motion of the
material. In a star, convection between adjacent layers is significant only when the temperature difference is great compared to the pressure difference. Convection occurs inside
the Sun within about 200,000 km of the surface. Convection is not a factor in the Sun’s
core.
(c) Radiation: the transfer of energy by electromagnetic waves. Inside the Sun and most
stars, radiation is the principal means of energy transport.
2. The material of the Sun is nearly opaque, so emitted radiation travels only a short distance (about 1 cm) before being absorbed and then reemitted. It takes a very long time for
a gamma-ray photon produced by a fusion reaction to reach the Sun’s surface. Thus any
information carried by such photons about the Sun’s core is lost.
3. Neutrinos are produced in the solar core and they rarely interact with matter. Their
very low mass implies a speed very close to the speed of light. Thus they can tell us about
the current conditions in the Sun’s core.
Solar Neutrinos and the Standard Solar Model
1. The standard solar model is the generally accepted theory of solar energy production.
2. The model predicts that so many neutrinos flow from the Sun that 65 billion neutrinos
pass through every cm2 of your body per second.
3. Neutrinos react very little with ordinary matter so detecting them is difficult. Astronomers want to measure the number of neutrinos reaching the Earth in order to check the
standard solar model.
4. Several large detectors have been built that confirm the solar origin of neutrinos but
report finding only 30% to 60% of the neutrinos predicted by theory. The discrepancy
between theory and experiment is referred to as the “solar neutrino problem,” and the
neutrino deficit depends on the energy of the neutrinos captured by the experiment.
5. Neutrinos come in three types, related to the electron and (its lesser known relatives)
the muon and tau particles. But most experiments were designed to detect electron neutrinos. According to the MSW theory, the observed neutrino deficit could be the result of
neutrino oscillations (neutrinos changing from one type to another during their flight
from the Sun to Earth).
6. The SNO experiment uses the properties of heavy water to detect all three types of
neutrinos. In 2002 it found that the number of electron neutrinos is about a third of the
total number of neutrinos reaching the detector, showing that neutrinos oscillate. The
SNO results are in excellent agreement with the predictions of the standard solar model.
7. Neutrino oscillations imply that neutrinos have mass, albeit very small. Since they are
the most numerous particles in the universe (other than photons), they contribute to a
small degree to its overall mass and influence its evolution.
10. So far we have only measured about 0.005% of the total number of neutrinos emitted
by the Sun. Future experiments will attempt to measure the remaining neutrinos and will
provide a further test for the standard solar model.
11-4 Helioseismology
1. In 1962 it was discovered that the Sun vibrates. The pulsations are caused by waves,
similar to sound waves, produced by pressure fluctuations in the turbulent convective
motions inside the Sun.
2. Helioseismology is the study of the propagation of these waves. It allows us to see
“through” the Sun and helps us predict solar activity.
Tools Of Astronomy: Solar Neutrino Experiments
1. The world’s first solar-neutrino detector began in the late 1960s and involved 378,000
liters of perchloroethylene, a dry-cleaning fluid containing chlorine, in an old gold mine
in Lead, South Dakota. When a solar neutrino strikes a chlorine atom with enough energy, a reaction occurs, transforming the chlorine into a radioactive isotope of argon. This
experiment measured about 30 of the numer of neutrinos predicted by the standard solar
model.
2. The Kamiokande experiment in the 1980s-1990s in Kamioka, Japan could also determine the direction of a neutrino’s travel, and confirmed that they do indeed come from
the Sun.
3. Those two detectors were sensitive only to high-energy neutrinos. Two detectors in
the 1990s, SAGE and GALLEX used the element gallium, and were able to detect the
low-energy neutrinos produced by the dominant reactions in the Sun’s core.
4. The Super-Kamiokande experiment, started in 1996, could distinguish different types
of neutrinos, and found 50% of the predicted amount coming from the Sun.
5. The SNO experiment in Sudbury, Canada, began in 1999 and used heavy water to detect all three types of neutrinos. This experiment confirmed the total number of neutrinos from the Sun was as expected but that they do change their type en-route to the Earth.
6. Support for the SNO results came in 2002 from the KamLAND project, which
measures anti-neutrinos from all nuclear plants in Japan. An antineutrino is a mirror image of a neutrino, and in this project we do not have to make any assumptions about the
source of these particles.
7. In the MINOS experiment, a beam of muon neutrinos travel underground from Fermilab to a detector in Soudan, MN; only 52% of them are detected, indicating neutrino
disappearance and therefore neutrino mass.
11-5 The Solar Atmosphere
1. The solar atmosphere is conveniently divided into three regions; the photosphere, the
chromosphere, and the corona.
The Photosphere
1. The photosphere is the visible “surface” of the Sun. It is the part of the solar atmosphere from which light is emitted into space. The photosphere is a very thin layer—400
km thick.
2. When observing the limb of the Sun (the apparent edge of the Sun as seen in the sky),
it appears darker than the center of the solar disk. That’s because when we observe the
limb of the Sun we see to a lesser depth because the line of sight is at a grazing angle.
3. The photosphere varies in temperature from about 6,500 K at its deepest to 4,400 K
near its outer edge. Overall, the light received from the photosphere is representative of
an object with temperature of about 5,800 K.
4. The pressure of the outer photosphere is 0.01 the pressure of the Earth’s surface.
Knowing the pressure and temperature we can calculate the density of particles in the
outer photosphere: it is about 0.0005 of the density of air at sea level on Earth.
5. The base of the photosphere shows granulation; the Sun’s surface is divided into
small convection cells. Granules are areas where hot material (light areas) is rising from
below and then descending (dark surroundings).
6. The chemical composition of the photosphere is (by mass) 78% hydrogen and 20%
helium; the remaining 2% consisting of 60 elements found on Earth. The similarity in
composition between Earth and Sun points to their common origins.
7. From our knowledge of nuclear fusion, we know the Sun’s core must hold more helium. Overall, the Sun (by mass) is theorized to be about 73% hydrogen, 25% helium, and
2% for all other elements.
The Chromosphere and Corona
1. The chromosphere is the region of the solar atmosphere (about 2,000 km thick) that
lies between the photosphere and the corona. It is not usually observable from Earth except during a total solar eclipse. It has a bright-line (emission) spectrum.
2. Corona is the outermost portion of the Sun’s atmosphere that can only be seen during
a total solar eclipse.
3. Narrow jets of gas (spicules) shoot upward from the chromosphere into the corona,
reaching a height of 6,000 to 10,000 km and lasting from 10 to 20 minutes.
4. The temperature increases as we move outward from the photosphere; it is as high as
30,000 K in the outer portions of the chromosphere and 2 million K in the corona (increasing rapidly in a transition region of about 300 km between the chromosphere and
corona).
5. The high temperatures in the chromosphere and especially the corona are not well understood. The leading hypothesis is that the heating results from the interaction between
the Sun’s magnetic field and its differential rotation. The presence of “solar moss” shows
that the field is highly organized in the transition region.
6. A prominence is an eruption of solar material beyond the disk of the Sun. Prominences may reach as high as a million kilometers above the photosphere. Slow-moving ones
can last for several days, while explosive ones can reach speeds of 1,500 km/s.
The Solar Wind
1. The solar wind is the continuous flow of nuclear particles (mostly protons and electrons) from the Sun.
2. Coronal holes, dark areas in an X-ray image of the Sun, correspond to regions of low
density where the magnetic field lines are open; they provide a corridor for charged particles to escape into space, generating the solar wind.
3. Due to the solar wind, the Sun loses about 6  1016 kg every year, a tiny fraction of its
total mass.
4. Near the Earth the solar wind travels at about 400 km/s and has a density of 2–10 particles/cm3.
5. Auroras form when the charged particles from the solar wind that have been trapped
by Earth’s magnetic field strike upper atmosphere molecules causing them to glow.
11-6 Sunspots and the Solar Activity Cycle
1. Dark spots on the Sun were first reported by the Chinese in the 5th century B.C. Galileo
and Thomas Harriott were the first Europeans to report sunspots in the early 17th century.
Sunspots are temporary phenomena lasting from a few hours to a few months.
2. Sunspots are about 1,500 K cooler than the surrounding photosphere. Thus they are
about 3 times less bright than their surrounding region.
3. The explanation for sunspots involves the Sun’s magnetic field. The strength of this
field can be measured using the Zeeman effect (the splitting of spectral lines by a strong
magnetic field).
4. Sunspots often appear in pairs, aligned in an east-west direction, and have opposite
magnetic polarities, one being north and the other south.
5. The magnetic field in a sunspot is about 1,000 times that of the surrounding photosphere.
6. In 1851, Schwabe discovered the sunspot cycle, which lasts about 11 years.
7. The Sun went through a period of inactivity during 1645–1715. This period corresponds to the “Little Ice Age” climatic period on Earth.
Modeling the Sunspot Cycle and the Sunspots
1. At a sunspot maximum, most spots occur about 35° north or south of the equator. As
the cycle progresses, the spots are seen closer and closer to the Sun’s equator. When the
spots reach the equator, the cycle is at a minimum and begins again.
2. The plot that shows the location and relative number of sunspots is called the butterfly
diagram.
3. The leading model for the sunspot cycle involves patterns of magnetic field lines, generated by the flow of the hot, ionized gases within the Sun’s interior, close to the boundary between the radiative and convective zones. At this boundary, the shear in the Sun’s
rotational velocity drives the formation of the field.
4. The lines form tubes that gradually become twisted due to the Sun’s differential rotation. Forced to the surface due to convection, these magnetic tubes become visible as
sunspots. Breaking the surface weakens the field lines and the sunspots die out.
5. The Sun’s overall magnetic field reverses from one sunspot cycle to the other. Thus the
entire magnetic cycle of the Sun is 22 years, twice the sunspot cycle of 11 years. During
each sunspot cycle, the polarity of the leading sunspot of a pair in a given hemisphere is
the same as the polarity of the Sun’s magnetic pole for that hemisphere.
6. Observations show that the Sun’s magnetic poles are connected to sunspots of opposite
polarity (trailing sunspots) near its equator by giant loops of hot plasma that extend into
the corona. The field steadily weakens due to the transfer of opposite magnetic flux to the
poles.
7. At the height of sunspot maximum, the field reverses and begins to grow in a new direction. During the flip the field is weak and uneven.
8. SOHO observations give us a better understanding of a sunspot region. The strong
magnetic field there behaves like a “plug” that stops the normal upward convective flow.
This cools the region above the plug, which then becomes denser and plunges downward
at high speeds, drawing the surrounding plasma and field toward the center of the sunspot. This increases the strength of the sunspot and this cycle repeats for as long as the
field is strong enough to behave like a plug.
Solar Flares and Coronal Mass Ejections
1. The turbulent magnetic field of the Sun causes explosive flare-ups called solar flares,
which occur mostly during sunspot maxima.
2. According to the sunspot model, flares occur when a great number of twisted tubes of
magnetic field lines release their energy at once through the photosphere.
3. Flares blast out large numbers of very energetic charged particles. In the case of the
largest flares, these particles can reach Earth in less than one hour.
4. Large solar flares cause spectacular auroras and can affect earthly radio transmissions
if the ionosphere is disrupted by the high-energy particles from the flare.
5. Coronal mass ejections (CMEs) are events in which hot coronal gas is suddenly
ejected into space at speeds of hundreds of km/s.
6. CMEs, flares and prominences are related phenomena.
7. Twisted field lines are seen in CMEs. Just like twisted coils of spring metal, they contain energy that is used to blast the material in space.