5.1: Exploration of the Sun

Exploring From a Distance

When scientists want to learn about something in nature they make observations.  Observations are easy when an object is close. For example, a botanist can study a plant in a greenhouse.  Sometimes a scientist needs to travel to make observations.  Geologists travel to many locations, sometimes very remote ones, to study geological structures such as mountains and volcanoes or to collect samples of rocks to bring back to a laboratory for further analysis. Astronomers study objects  that are very far away by collecting the light given off by stars. Other objects may not give off their own visible light, but reflect light from as the Moon reflects the light of the Sun. Astronomers don't have to rely on visible light because the entire electromagnetic spectrum can be used to gain information.  This section is about the tools that help astronomers learn about objects in the Solar System and beyond. First, we need to review observing safety.

Observing the Sun Safely

Because looking at the Sun for even a brief time can cause permanent eye damage, never look directly at the Sun without proper safety equipment.

When observing the Sun, you MUST use:

• A proper solar filter
  • Position the filter in front of your eyes, a telescope or binoculars.
  • A solar filter is safe because it does not transmit ultraviolet or infrared radiation, both of which are much more harmful to your eyes than visible light.
  • The filter also decreases the Sun’s brightness to a comfortable level.

or

• A pinhole camera can be used to safely – and indirectly – view the sun and sunspots.
• You can make your own pinhole camera using materials found at home.

Seeing Spots on the Sun

The Sun's surface has small, darker areas called sunspots. There are records going back over a thousand years from observers who saw sunspots when haze or mist reduced the Sun’s brightness. The earliest recorded sightings of sunspots were made by several astronomers including Thomas Harriot, Galileo Galilei, and Johannes Fabricius. After his first observation in 1611, Fabricius used a camera obscura and telescope together, projecting the image of the Sun onto a blank surface. This method both protected his eye sight and made it possible to observe that sunspots move across the surface of the Sun. 

 

Solar Rotation

Like storms on Earth, drift slowly across the surface of the Sun. In 1612, Galileo recorded the apparent motions of sunspots as the turning Sun carried them across its disk. From these observations, he was able to estimate that the Sun rotates on its axis with a rotation period of about 1 month. By (Figure 15.14), . Our star turns in a west-to-east direction, like the orbital motions of the planets. The Sun is made of hot, ionized gas, called plasma, and does not rotate rigidly, the way a solid body like Earth does. Modern observations show that the speed of rotation of the Sun varies according to latitude, that is, it’s different as you go north or south of the Sun’s equator. The rotation period is about 25 days at the equator, 28 days at latitude 40°, and 36 days at latitude 80°. We call this behavior differential rotation, which you will see again in the chapters about the outer planets.

 

Figure 15.14 : Sunspot Movement. This sequence of photographs of the Sun’s surface tracks the movement of sunspots across the visible hemisphere of the Sun. On March 30, 2001, this group of sunspots extended across an area about 13 times the diameter of Earth. This region produced many flares and coronal mass ejections. (Openstax Astronomy)

Weighing the Sun

Star characteristics consider physical characteristics such as stellar mass, size, surface temperature, and luminosity. These basic characteristics assist in detailing specifics about a star; not all stars are the same and these variations in characteristics can be significant. The distance from the Sun and Earth to a specific star is also important as astronomers detail the star’s characteristics. Of these characteristics, stellar mass is the most important characteristic. The mass of a star determines things, such as how quickly a star will consume its stellar fuel through the fusion process to the star’s final ending when it has used all of its fuel. Yet how does one measure a star’s mass? How do you weigh a star?

It turns out that Isaac Newton studied this question. He turned to Kepler’s Third Law, where one measures the period and average distance of the object’s orbit about a star. Yet, you need two objects – a star and an object orbiting a star – to use this solution. It turns out that over 50% of all stars have a companion star. So, astronomers use this adaptation on Kepler’s Third Law developed by Newton to measure the mass of the two binary stars.

Kepler’s 3rd Law — a ³ = kP ²

Where a is the orbiting object’s semi major axis, P is the orbiting object’s period to orbit, and k is a constant, referred to as Kepler’s constant.

By examining the color of each of the stars in the binary system, you can compare two single stars with the same colors.

Stellar mass is usually related to the mass of the Sun, where the Sun equals 1 m _sun , 1 m _⊙ , or 1 solar mass. The bright star Sirius, the Dog Star in the constellation Canis Major, is about 2. 02 m _sun. One of the most massive stars is Eta Carinae, with a mass somewhere between 100 to 150 times the mass of the Sun, 100-150 m _⊙ .

Stellar mass units — m _sun :: m _⊙ :: solar mass

A star’s mass will vary over its lifetime, depending if it adds, or accretes, mass from another star, loses mass to another star, or simply loses mass through the normal processes, such as through its stellar wind or pulsating outputs.

Stars are occasionally classed by their stellar masses based upon their evolutionary behavior as stars approach the end of their nuclear fusion. In the next module, we will introduce the classes of stars, based on their solar masses.

Stellar size refers to a star’s diameter or radius. Stars range in diameter, from neutron stars with diameters of about 40 kilometers or 25 miles, to supergiants with diameters of approximately 900,000,000 kilometers or 540,000,000 miles — about 650 times the Sun’s diameter.

A star’s surface temperature, measured in Kelvin, K, is dependent on the star’s diameter and the rate of energy production at the stellar core, and is measured at the star’s photosphere. An estimate of the surface temperature is the star’s color, often called the color index. Annie Jump Cannon was the first to sort spectral data and designed the stellar spectral classes.

The hotter the star, the whiter it will appear, whereas the cooler the star, the redder it will appear. Think of heating a piece of metal; the hotter it is, the whiter the metal will appear. As the metal cools, it will appear orange and then red in color. The reddish-colored metal is still hot, yet cooler than when the metal was white blue-white in color.

Stellar luminosity is the amount of light and other radiant energy released by a star. A star’s luminosity is dependent on its diameter (sometimes noted as the star’s radius, d = 2r) and its surface temperature.

Taking the Sun's Temperature

Can you imagine an astronomer trying to stick a thermometer in the Sun? Not only is that an extremely dangerous thing to do, the thermometer would be destroyed long before it touched the Sun.  How do astronomers know how hot the Sun is?  The next sections will demonstrate how light and temperature are connected to each other.

 Some astronomical objects emit mostly infrared radiation, others mostly visible light, and still others mostly ultraviolet radiation. What determines the type of electromagnetic radiation emitted by the Sun, stars, and planets? The answer often turns out to be their temperature.

Temperature is operationally defined to be what we measure with a thermometer.  discussed in physics section - local summary here The temperature of something is a measure of the average thermal energy of the particles that make up that object.

Summarize Blackbody from light 

In the section below, we will review the particles that make up matter, known as atoms. For now, you need to know that atoms are made up three different types of particles, one of which is called the electron. Electrons all have a negative charge. In the section about light, we learned that when a charged particle moves, it releases electromagnetic radiation, 

This motion at the microscopic level is responsible for much of the electromagnetic radiation on Earth and in the universe. As atoms and molecules move about and collide, or vibrate in place, their electrons give off electromagnetic radiation. The characteristics of this radiation are determined by the temperature of those atoms and molecules. In a hot material, for example, the individual particles vibrate in place or move rapidly from collisions, so the emitted waves are, on average, more energetic. And recall that higher energy waves have a higher frequency. In very cool material, the particles have low-energy atomic and molecular motions and thus generate lower-energy waves.

To understand, in more quantitative detail, the relationship between temperature and electromagnetic radiation, we imagine an idealized object called a blackbody. Such an object (unlike your sweater or your astronomy instructor’s head) does not reflect or scatter any radiation, but absorbs all the electromagnetic energy that falls onto it. The energy that is absorbed causes the atoms and molecules in it to vibrate or move around at increasing speeds. As it gets hotter, this object will radiate electromagnetic waves until absorption and radiation are in balance. We want to discuss such an idealized object because, as you will see, stars behave in very nearly the same way.

The radiation from a blackbody has several characteristics, as illustrated in Figure 5.8. The graph shows the power emitted at each wavelength by objects of different temperatures. In science, the word power means the energy coming off per second (and it is typically measured in watts, which you are probably familiar with from buying lightbulbs).

Figure 5.8 : Light and Temperature This graph shows in arbitrary units how many photons are given off at each wavelength for objects at four different temperatures. The wavelengths corresponding to visible light are shown by the colored bands. Note that at hotter temperatures, more energy (in the form of photons) is emitted at all wavelengths. The higher the temperature, the shorter the wavelength at which the peak amount of energy is radiated (Openstax Astronomy).

First of all, notice that the curves show that, at each temperature, our blackbody object emits radiation (photons) at all wavelengths (all colors). This is because in any solid or denser gas, some molecules or atoms vibrate or move between collisions slower than average and some move faster than average. So when we look at the electromagnetic waves emitted, we find a broad range, or spectrum, of energies and wavelengths. More energy is emitted at the average vibration or motion rate (the highest part of each curve), but if we have a large number of atoms or molecules, some energy will be detected at each wavelength.

Second, note that an object at a higher temperature emits more power at all wavelengths than does a cooler one. In a hot gas (the taller curves in Figure 5.8), for example, the atoms have more collisions and give off more energy. In the real world of stars, this means that hotter stars give off more energy at every wavelength than do cooler stars.

Third, the graph shows us that the higher the temperature, the shorter the wavelength at which the maximum power is emitted. Remember that a shorter wavelength means a higher frequency and energy. It makes sense, then, that hot objects give off a larger fraction of their energy at shorter wavelengths (higher energies) than do cool objects. You may have observed examples of this rule in everyday life. When a burner on an electric stove is turned on low, it emits only heat, which is infrared radiation, but does not glow with visible light. If the burner is set to a higher temperature, it starts to glow a dull red. At a still-higher setting, it glows a brighter orange-red (shorter wavelength). At even higher temperatures, which cannot be reached with ordinary stoves, metal can appear brilliant yellow or even blue-white.

We can use these ideas to come up with a rough sort of “thermometer” for measuring the temperatures of stars. Because many stars give off most of their energy in visible light, the color of light that dominates a star’s appearance is a rough indicator of its temperature. If one star looks red and another looks blue, which one has the higher temperature? Because blue is the shorter-wavelength color, it is the sign of a hotter star. (Note that the temperatures we associate with different colors in science are not the same as the ones artists use. In art, red is often called a “hot” color and blue a “cool” color. Likewise, we commonly see red on faucet or air conditioning controls to indicate hot temperatures and blue to indicate cold temperatures. Although these are common uses to us in daily life, in nature, it’s the other way around.)

We can develop a more precise star thermometer by measuring how much energy a star gives off at each wavelength and by constructing diagrams like Figure 5.8. The location of the peak (or maximum) in the power curve of each star can tell us its temperature. The average temperature at the surface of the Sun, which is where the radiation that we see is emitted, turns out to be 5800 K. (Throughout this text, we use the kelvin or absolute temperature scale. On this scale, water freezes at 273 K and boils at 373 K. All molecular motion ceases at 0 K. The various temperature scales are described in Appendix D.) There are stars cooler than the Sun and stars hotter than the Sun.

The wavelength at which maximum power is emitted can be calculated according to the equation

$${\text{λ}}_{\text{max}}=$$

where the wavelength is in nanometers (one billionth of a meter) and the temperature is in K (the constant $3\times 10^6\text{has units of}$=3×106nm Kλmax=3×106nm K1200 nm=2500K

Since this star has a peak wavelength that is at a shorter wavelength (in the ultraviolet part of the spectrum) than that of our Sun (in the visible part of the spectrum), it should come as no surprise that its surface temperature is much hotter than our Sun’s.

We can also describe our observation that hotter objects radiate more power at all wavelengths in a mathematical form. If we sum up the contributions from all parts of the electromagnetic spectrum, we obtain the total energy emitted by a blackbody. What we usually measure from a large object like a star is the energy flux, the power emitted per square meter. The word flux means “flow” here: we are interested in the flow of power into an area (like the area of a telescope mirror). It turns out that the energy flux from a blackbody at temperature T is proportional to the fourth power of its absolute temperature. This relationship is known as the Stefan-Boltzmann law and can be written in the form of an equation as

$$F=$$

where F stands for the energy flux (in units of watts per square meter), T is given in Kelvins, and σ (Greek letter sigma) is a constant number $(5.67\times {10}^{-8})$.

Notice how impressive this result is. Increasing the temperature of a star would have a tremendous effect on the power it radiates. If the Sun, for example, were twice as hot—that is, if it had a temperature of 11,600 K—it would radiate 2⁴, or 16 times more power than it does now. Tripling the temperature would raise the power output 81 times. Hot stars really shine away a tremendous amount of energy.

The history of quantum mechanics starts with the study of thermal physics. With the development of theory of electromagnetism and thermodynamics in the 18th century, physicists felt confident enough to attempt to develop a theory of blackbody radiation, which describes the EM spectrum radiated by a hot solid. It is called blackbody radiation, because an ideal radiator is jet black. It absorbs all radiation incident on it, and thus it can be at thermal equilibrium with the surrounding while emitting more radiation than a white object, which reflects radiation. So such a hypothetical object is called a blackbody. The first, classical attempt at theoretical description of this phenomenon was nothing short of a catastrophe. Experimentally observed spectrum showed less intensity of EM radiation at both the very short wavelength limit and the very long wavelength limit, but the classical theory prediction predicted ever-increasing intensity at shorter and shorter wavelengths. (See Figure \(\PageIndex{1}\).)

 

This result came to be called ultraviolet catastrophe, due to the catastrophically large intensity of radiation predicted in the ultraviolet range of EM spectrum. This problem stumped many physicists until Max Planck came up with a novel suggestion.

The spectrum of blackbody radiation was extensively studied and well known. The German physicist Max Planck (1858–1947) first guessed a functional form of intensity dependence on the wavelength of EM radiation, which is now known as the Planck Law. The Planck Law was a good fit to the experimental results tapering off both at the long-wavelength and the short-wavelength limit, but it was only a phenomenological description that did not include a satisfying explanation of why it was so.

  

This led Planck to look for a reason why intensity of radiation should decrease at the short-wavelength, or high-frequency limit. Using the idea that atoms and molecules in a body act like oscillators to absorb and emit radiation, he guessed that the energy of these oscillators are quantized. That is, the energy of the oscillators could change only by a discrete amount, and this discrete amount of energy change was \(\Delta E=h f\), where \(f\) is frequency of the oscillator and \(h\) is a fundamental constant of nature that we now call Planck's constant, given by

\[h=6.626 \times 10^{-34} \mathrm{~J} \cdot \mathrm{s}. \nonumber \]

Without getting into the detailed mathematical derivation, we can see intuitively why this assumption would lead to the result that Planck was looking for. An object at thermal equilibrium has a certain amount of thermal energy, which is associated with kinetic energy of microscopic movements and vibrations. If this energy is randomly distributed, both very high-frequency oscillators and very low-frequency oscillators would have the same thermal energy on average. But if it was true that this thermal energy is quantized, and it has to come in integer units of \(hf\), then at very high frequencies, there is not going to be enough thermal energy for even one unit of \(hf\). This places an upper limit on maximum possible frequency of these oscillators, for a given amount of thermal energy per oscillator, resulting in the tapering behavior for intensity of radiation emitted by these oscillators at the higher frequency (or shorter wavelengths).

This model would also explain why the peak wavelength shifts to the shorter wavelength (higher frequency) at higher temperatures. At higher temperatures, there is more thermal energy available per oscillator, so the upper limit on maximum possible frequency will be higher, since larger unit size \(hf\) is allowed.

Note that Planck’s constant \(h\) is a very small number. So for an infrared frequency of \(10^{14} \mathrm{~Hz}\) being emitted by a blackbody, for example, the difference between energy levels is only \(\Delta E=h f=\left(6.63 \times 10^{-34} \mathrm{~J} \cdot \mathrm{s}\right)\left(10^{14} \mathrm{~Hz}\right)=6.63 \times 10^{-20} \mathrm{~J}\) or about 0.4 eV. This 0.4 eV of energy is significant compared with typical atomic energies, which are on the order of an electron volt, or thermal energies, which are typically fractions of an electron volt. But on a macroscopic or classical scale, energies are typically on the order of joules. Even though the macroscopic energies are quantized, the quantum steps are too small to be noticed. This is an example of the correspondence principle. For a large object, quantum mechanics produces results indistinguishable from those of classical physics.

Remote Samples of the Sun

summarize formation of spectral lines from light

The Value of Stellar Spectra

When Newton described the laws of refraction and dispersion in optics, and observed the solar spectrum, all he could see was a continuous band of colors. If the spectrum of the white light from the Sun and stars were simply a continuous rainbow of colors, astronomers would have little interest in the detailed study of a star’s spectrum once they had learned its average surface temperature. In 1802, however, William Wollaston built an improved spectrometer that included a lens to focus the Sun’s spectrum on a screen. With this device, Wollaston saw that the colors were not spread out uniformly, but instead, some ranges of color were missing, appearing as dark bands in the solar spectrum. He mistakenly attributed these lines to natural boundaries between the colors. In 1815, German physicist Joseph Fraunhofer, upon a more careful examination of the solar spectrum, found about 600 such dark lines (missing colors), which led scientists to rule out the boundary hypothesis (Figure 5.11).

Figure 5.11 : Visible Spectrum of the Sun. Our star’s spectrum is crossed by dark lines produced by atoms in the solar atmosphere that absorb light at certain wavelengths. (Openstax Astronomy)

Later, researchers found that similar dark lines could be produced in the spectra (“spectra” is the plural of “spectrum”) of artificial light sources. They did this by passing their light through various apparently transparent substances—usually containers with just a bit of thin gas in them.

These gases turned out not to be transparent at all colors: they were quite opaque at a few sharply defined wavelengths. Something in each gas had to be absorbing just a few colors of light and no others. All gases did this, but each different element absorbed a different set of colors and thus showed different dark lines. If the gas in a container consisted of two elements, then light passing through it was missing the colors (showing dark lines) for both of the elements. So it became clear that certain lines in the spectrum “go with” certain elements. This discovery was one of the most important steps forward in the history of astronomy.

What would happen if there were no continuous spectrum for our gases to remove light from? What if, instead, we heated the same thin gases until they were hot enough to glow with their own light? When the gases were heated, a spectrometer revealed no continuous spectrum, but several separate bright lines. That is, these hot gases emitted light only at certain specific wavelengths or colors.

When the gas was pure hydrogen, it would emit one pattern of colors; when it was pure sodium, it would emit a different pattern. A mixture of hydrogen and sodium emitted both sets of spectral lines. The colors the gases emitted when they were heated were the very same colors as those they had absorbed when a continuous source of light was behind them. From such experiments, scientists began to see that different substances showed distinctive spectral signatures by which their presence could be detected  Just as your signature allows the bank to identify you, the unique pattern of colors for each type of atom (its spectrum) can help us identify which element or elements are in a gas.

Figure 5.12 : Emission Spectra of Elements. Each type of glowing gas (each element) produces its own unique pattern of lines, so the composition of a gas can be identified by its spectrum. The spectra of sodium, hydrogen, calcium, and mercury gases are shown here (Openstax Astronomy).

Composition of the Sun

The Sun, like all stars, is an enormous ball of extremely hot, largely ionized gas, shining under its own power. And we do mean enormous. The Sun could fit 109 Earths side-by-side across its diameter, and it has enough volume (takes up enough space) to hold about 1.3 million Earths.

The Sun does not have a solid surface or continents like Earth, nor does it have a solid core. However, it does have a lot of structure and can be discussed as a series of layers, not unlike an onion. In this section, we describe the huge changes that occur in the Sun’s extensive interior and atmosphere, and the dynamic and violent eruptions that occur daily in its outer layers.

 

The fact that our Sun and the stars all have similar compositions and are made up of mostly hydrogen and helium was first shown in a brilliant thesis in 1925 by Cecilia Payne-Gaposchkin, the first woman to get a PhD in astronomy in the United States (Figure 15.3). However, the idea that the simplest light gases—hydrogen and helium—were the most abundant elements in stars was so unexpected and so shocking that she was persuaded her analysis of the data must be wrong. At the time, she wrote, “The enormous abundance derived for these elements in the stellar atmosphere is almost certainly not real.” Even scientists sometimes find it hard to accept new ideas that do not agree with what everyone “knows” to be right.

  

Before Payne-Gaposchkin’s work, everyone assumed that the composition of the Sun and stars would be much like that of Earth. It was 3 years after her thesis that other studies proved beyond a doubt that the enormous abundance of hydrogen and helium in the Sun is indeed real. (And, as we will see, the composition of the Sun and the stars is much more typical of the makeup of the universe than the odd concentration of heavier elements that characterizes our planet.)

Most of the elements found in the Sun are in the form of atoms, with a small number of molecules, all in the form of gases: the Sun is so hot that no matter can survive as a liquid or a solid. In fact, the Sun is so hot that many of the atoms in it are ionized, that is, stripped of one or more of their electrons. This removal of electrons from their atoms means that there is a large quantity of free electrons and positively charged ions in the Sun, making it an electrically charged environment—quite different from the neutral one in which you are reading this text. (Scientists call such a hot ionized gas a plasma.)

In the nineteenth century, scientists observed a spectral line at 530.3 nanometers in the Sun’s outer atmosphere, called the corona (a layer we will discuss in a minute.) This line had never been seen before, and so it was assumed that this line was the result of a new element found in the corona, quickly named coronium. It was not until 60 years later that astronomers discovered that this emission was in fact due to highly ionized iron—iron with 13 of its electrons stripped off. This is how we first discovered that the Sun’s atmosphere had a temperature of more than a million degrees.

Let’s begin by asking what the solar atmosphere is made of. As explained in Radiation and Spectra, we can use a star’s absorption line spectrum to determine what elements are present. It turns out that the Sun contains the same elements as Earth but not in the same proportions. About 73% of the Sun’s mass is hydrogen, and another 25% is helium. All the other chemical elements (including those we know and love in our own bodies, such as carbon, oxygen, and nitrogen) make up only 2% of our star. The 10 most abundant gases in the Sun’s visible surface layer are listed in Table 15.2. Examine that table and notice that the composition of the Sun’s outer layer is very different from Earth’s crust, where we live. (In our planet’s crust, the three most abundant elements are oxygen, silicon, and aluminum.) Although not like our planet’s, the makeup of the Sun is quite typical of stars in general.   

Table 15.2: The Abundance of Elements in the Sun

Element	Percentage by Number of Atoms	Percentage By Mass
Hydrogen	92.0	73.4
Helium	7.8	25.0
Carbon	0.02	0.20
Nitrogen	0.008	0.09
Oxygen	0.06	0.80
Neon	0.01	0.16
Magnesium	0.003	0.06
Silicon	0.004	0.09
Sulfur	0.002	0.05
Iron	0.003	0.14

Spacecraft 

As of July 14, 2025, there have been 17 spacecraft (not including the ones in Earth orbit) that have successfully studied the Sun and two spacecraft that have yet to begin their mission. These spacecraft have returned a large amount of data that astronomers have used to understand how the Sun works.

High-Resolution Images

• Define the word resolution in relation to digital images
• Describe how methods and limitations on increasing image resolution

need to find sources, images showing high & low resolution 

False-Color Images

Attributions

This section was adapted from   and published under a Creative Common Attribution License 4.0 and from "10.8: Studying the Sun" and "10.10: The National Solar Observatory" in Introduction to Astronomy (Lumen) originally written by Lumen Learning, and published under a Creative Common Attribution License 4.0.

https://phys.libretexts.org/Bookshel...haracteristics

https://phys.libretexts.org/Bookshel...body_Radiation

References

• List of Solar System probes. (2025, Nov 14). In Wikipedia.
• Image Resolution. (2025, June 6). In Wikipedia.
• False color. (2025, Oct 5). In Wikipedia.
• https://en.wikipedia.org/wiki/Johannes_Fabricius

Further Exploration

• Spectrum Constructor by Foothill Astrosims CC BY NC SA
• Build an Atom by PhET 
• Hydrogen Atom simulation from Foothill Astrosims
• NASA Sun:Exploration