3. Atomic and Subatomic Particles


Fig3_1 The helium atom

Fig3_1 The helium atom

Fig. 3.1 . An atom of helium contains two electrons that swarm about the atomís nuclear center in a cloud of largely empty space. The shading shows that the electrons can be anywhere, but are most likely to be found near the center of the atom. The magnified nucleus of the helium atom consists of two protons and two neutrons bound together by a strong nuclear force. The nucleus and each of its four particles are spherically symmetric. The size of the helium nucleus is about 1 fermi, or 1 fm, which is equivalent to 10-15 m. The atom is about 100,000 times bigger than the nucleus, with an atom size of about 105 fm or 10-10 m.


Fig3_2 Random and ordered motion

Fig3_2 Random and ordered motion

Fig. 3.2 . Particles within a hot gas (left) move here and there in random directions that continually change as the result of collisions between particles. This supports the gas against inward gravitational forces. A planet or star (right) moves along a well defined, ordered trajectory determined by external gravitational forces upon it. When a large number of stars have gathered together and are confined within a star cluster, the stars also move in random directions, supporting their combined gravitational pull.


Fig3_3 Maxwell Distribution of Particle Speeds

Fig3_3 Maxwell Distribution of Particle Speeds

Fig. 3.3 . The speeds of particles with the same mass and three different temperatures. The peak of this distribution shifts to higher speeds at higher temperatures. There is a small fraction of particles having high speed, residing in the high-speed tail of the distribution, and this fraction increases with temperature. The fraction of particles with low speed becomes smaller at higher temperatures but does not vanish. The peak also shifts to higher speeds at lower mass when the temperature is unchanged. The Scottish scientist James Clerk Maxwell (1831?1879) derived this distribution in 1873.


Fig3_4 Earth's layered atmosphere

Fig3_4 Earth

Fig. 3.4 . The pressure of our atmosphere (right scale) decreases with altitude (left scale). This is because fewer particles are able overcome the Earthís gravitational pull and reach higher altitudes. The temperature (bottom scale) also decreases steadily with height in the ground-hugging troposphere, but the temperature increases in two higher regions that are heated by the Sun. They are the stratosphere, with its critical ozone layer, and the ionosphere. The stratosphere is mainly heated by ultraviolet radiation from the Sun, and the ionosphere is created and modulated by the Sunís x-ray and extreme ultraviolet radiation that breaks the atmospheric molecules apart, and strips electrons off their component atoms to produce ions. The process of ionization by the Sunís invisible rays releases heat to warm the ionosphere, so the temperature rises with altitude in it. In the ionosphere, at about 100 km to 500 km above the ground, the temperatures skyrocket to higher values than anywhere else in the entire atmosphere. At higher altitudes, the atmosphere thins out into the exosphere, or the ďexit to the outside sphere.Ē The temperature is so hot out there, and the particles move so fast, that some atoms and molecules may slowly evaporate away.


Fig3_5 States of Matter

Fig3_5 States of Matter

Fig. 3.5 . The locations of atoms, large dashed circles, and their component electrons, small filled circles, and central nucleus, large filled circles, for the gaseous (left), liquid or solid (center) and plasma (right) states of matter. In the gaseous state, the atoms are widely separated and free to move about. The atoms are practically touching each other in the solid, and liquid, state. At sufficiently high temperature and pressure, the atoms cease to exist, and the plasma state is created. The atoms are torn into their constituents by frequent collisions at the high temperatures. The plasma consists of bare nuclei and unattached electrons moving about in random directions within the former empty space of atoms. In the plasma state, matter regains the compressibility of the gaseous state, and plasma behaves like a gas.


Fig3_6 Internal compression of the Sun

Fig3_6 Internal compression of the Sun

Fig. 3.6 . The variation of pressure, luminosity, temperature and mass density with fractional radial distance from the Sun's center (left) to its visible disk (right). At the Sun's center, the temperature is 15.6 million K and the mass density is 151,300 kg m-3; the central pressure is 2.33 x 1016 Pa or 233 billion times that of the Earth's atmosphere at sea level (one bar equivalent to 100,000 Pa). Nuclear reactions occur only in the central core, out to about 25 percent of the Sunís radius. The energy produced in the core is transported by radiation out to 71 percent of the starís radius, where the temperature has dropped to about 2 million K and the density has fallen to about 200 kg m-3. The energy is then transported by convection out to the Sunís visible disk, known as the photosphere, where the temperature is 5780 K, and the pressure and density have dropped off the scales of the graph.


Fig3_7 Spectroheliograph

Fig3_7 Spectroheliograph

Fig. 3.7 . A small section of the Sunís image at the focal plane of a telescope is selected with a narrow entry slit, S1, and this light passes to a diffraction grating to produce a spectrum. A second slit, S2, at the focal plane selects a specific wavelength from the spectrum. If the plate containing the two slits is moved horizontally, then the entrance slit passes adjacent strips of the solar image. The light leaving the moving exit slit then builds up an image of the Sun at a specific wavelength.


Fig3_8 Absorption and Emission Lines

Fig3_8 Absorption and Emission Lines

Fig. 3.8 . The spectrum of a star or other cosmic object displays the intensity of its radiation as a function of wavelength, denoted by the Greek symbol ?. Any hot gas radiates at all wavelengths, producing a continuum spectrum with an intensity that varies with the wavelength and depends on the temperature. When this thermal radiation passes through an outer, cooler layer of a star, the atoms in this layer can produce absorption at a specific wavelength. This feature is called an absorption line because it looks like a line in the spectrum. When atoms are excited at high temperatures, they can radiate an emission line. The line wavelength tells us which atom is responsible for the absorption or emission. The intensity of a stellar line is related to both the number of atoms and the physical conditions in the starís atmosphere. The lines have a width and a shape that can be observed under close scrutiny in wavelength, and they provide information about the local physical conditions. The motion of the absorbing or emitting atoms, for example, broadens the lines (see Fig. 3.13).


Fig3_9 Element abundance in solar photosphere

Fig3_9 Element abundance in solar photosphere

Fig. 3.9 . The relative abundance of the elements in the Sunís visible disk, the photosphere, plotted as a function of atomic number. The atomic number, denoted by Z, is the number of protons in the atomís nucleus, or roughly half the atomic weight. Heavy elements, with high atomic number, are less abundant than light ones, with low atomic number, and the most abundant element in the Sun is the lightest element, hydrogen. Helium is the second most abundant solar element. The abundance data are plotted in a logarithmic scale normalized to a million million, or 1.0 x 1012, for hydrogen. [Adapted from Martin Asplund, Nicolas Grevesse, A. Jacques Sauval and Pat Scott, The Chemical Composition of the Sun, Annual Review of Astronomy and Astrophysics 47, 481-522 (2009).]


Fig3_10 Balmer line of hydrogen

Fig3_10 Balmer line of hydrogen

Fig. 3.10 . The spectrum of the Sunís optically visible light exhibits four strong absorption lines that are attributed to hydrogen whose line wavelengths are spaced closer together at shorter wavelengths (top). These lines are designated H? at a red wavelength of 656.3 nm, H? at a wavelength of 486.1 nm, H? at the blue 434.1 nm and H? at the violet 410.2 nm, where one nanometer = 1 nm = 10-9 m. These spectral features originate when an electron in a hydrogen atom moves from a low to high electron orbit, whose orbital energy is a function of the integer n (bottom). They have been named the Balmer lines after the Swiss mathematics teacher Johann Balmer (1825?1898) who first derived an equation in 1895 that describes their wavelengths in terms of integers.


Fig3_11 Bohr atom

Fig3_11 Bohr atom

Fig. 3.11 . In this model, proposed in 1913 by the Danish physicist Niels Bohr (1885-1962), a hydrogen atomís one electron revolves around the hydrogen nucleus, a single proton, in well-defined orbits described by the integer n = 1, 2, 3, 4, 5, Ö An electron absorbs or emits radiation when it makes a transition between these allowed orbits. The electron can jump upward, to orbits with larger n, by absorption of a radiation photon of exactly the right energy, equal to the energy difference between the orbits; the electron can jump down to lower orbits, of smaller n, with the emission of radiation of that same energy and wavelength. Transitions that begin or end on the n = 2 orbit define the Balmer series that is observed at visible wavelengths. They are designated by Ha, Hb, Hg, .... The Lyman series, with transitions from the first orbit at n = 1, is detected at ultraviolet wavelengths. The orbits are not drawn to scale for the size of their radius increases with the square of the integer n.


Fig3_12 Doppler effect

Fig3_12 Doppler effect

Fig. 3.12 . A stationary source of radiation (top) emits regularly spaced light waves that get stretched out or scrunched up if the source moves (bottom). Here we show a star moving away (bottom right) from the observer (bottom left). The stretching of light waves that occurs when the source moves away from an observer along the line of sight is called a redshift, because red light waves are relatively long visible light waves; the compression of light waves that occurs when the source moves along the line of sight toward an observer is called a blueshift, because blue light waves are relatively short. The wavelength change, from the stationary to moving condition, is called the Doppler shift, and its size provides a measurement of radial velocity, or the speed of the component of the sourceís motion along the line of sight. The Doppler effect is named after the Austrian physicist Christian Doppler (1803-1853), who first considered it in 1842.


Fig3_13 Effects that broaden spectral line

Fig3_13 Effects that broaden spectral line

Fig. 3.13 . The motion of absorbing or emitting atoms can broaden a line to the short-wavelength and long-wavelength sides of the rest, or non-moving, wavelength, here denoted by lL (top). The Doppler effect describes the broadening. When the motion is due to the heat, or temperature, of the radiating atoms, the effect produces thermal broadening, and when the average temperature of a collection of atoms increases the thermal broadening becomes wider. An intense magnetic field can split a single line at wavelength lL into two components by the Zeeman effect (bottom). The wavelength difference DlL between the split lines is proportional to the strength of the magnetic field.