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Heliophysics and Stellar Atmospheres > Semester 2 - Definitions > Flashcards

Semester 2 - Definitions Flashcards

1
Q

What are the layers of the solar atmosphere?

A

Photosphere, Chromosphere, Transition Region, Corona.

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2
Q

What defines the solar surface in the photosphere?

A

The location where solar gas is opaque at a wavelength of 500 nm.

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3
Q

What are the typical plasma temperatures in different layers of the solar atmosphere?

A

Photosphere: 5800K
Chromosphere: 4000K to 10^4K
Transition Region: 10^4K to 10^6K
Corona: 1-2MK

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4
Q

Why does the solar disc have a sharp edge?

A

At the solar edge the optical thickness adrops to τ &laquo_space;1. This is at the photosphere where the opacity is large due to Thomson scattering.

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5
Q

What is limb darkening?

A

Limb darkening occurs because the optical thickness is greater along the limb of the Sun compared to its center, causing photons to scatter away or be absorbed along the line of sight.

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6
Q

What is Thomson scattering, and why is it significant in the outer layers of the Sun?

A

Thomson scattering occurs when a photon is scattered off a free electron without changing its energy. It is significant in the outer layers of the Sun due to the large opacity caused by this scattering process.

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7
Q

Why does the solar photosphere have a sharp edge?

A

The exponential drop in intensity with height due to the decrease in electron density leads to a sharp transition in opacity.

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8
Q

What are the most abundant elements in the solar photosphere?

A

Hydrogen and Helium.

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9
Q

What is the significance of minor ions in the solar atmosphere?

A

Minor ions in the solar atmosphere are useful for plasma diagnostics

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10
Q

What is the radiative instability hypothesis?

A

It explains the sudden increase in temperature seen in the transition region of the solar atmosphere, suggesting a balance between heating and cooling rates.

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11
Q

What role does thermal conduction play in the transition region and corona?

A

Thermal conduction transports energy backwards towards the lower chromosphere in the transition region and contributes to the radial increase in temperature observed in the corona.

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12
Q

How does Kirchoff’s laws explain the presence of absorption lines.

A

Kirchoff’s laws predict that absorption lines will be observed due to the presence of cooler material higher up in the atmosphere.

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13
Q

What is the chromosphere’s optical property, and how is it observed?

A

The chromosphere is optically thin in continuum regions but optically thick for bright lines such Hα.

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14
Q

How is the solar chromosphere revealed?

A

by pink/redish emission in eclipse images.

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15
Q

How is the chromosphere connected to the corona, and what lies between them?

A

The chromosphere connects the photosphere to the corona, with a transition region between them that is approximately 100km thick.

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16
Q

What are the main heating mechanisms proposed for the chromosphere and corona?

A

Acoustic (sound) waves and Magnetic Waves.

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17
Q

How do acoustic waves contribute to heating in the chromosphere and corona?

A

Acoustic waves, generated by photospheric motions, propagate upwards and can heat the chromosphere. However, they break too early or are reflected back towards the chromosphere, limiting their ability to explain high coronal temperatures.

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18
Q

What are Alfven waves, and how do they contribute to heating in the corona?

A

Alfven waves are magnetic fluctuations that propagate along magnetic field lines into the corona. They damp energy and contribute to the temperature structure of the corona.

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19
Q

How do Alfven waves differ from sound waves?

A

Unlike sound waves, Alfven waves are incompressible magnetic fluctuations, balancing magnetic tension with plasma inertia.

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20
Q

Where are Alfven waves observed, and how are they believed to dissipate their energy?

A

Alfven waves are observed in the solar wind and are believed to dissipate their energy in the solar wind via turbulent cascades to smaller scales.

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21
Q

What are the two main mechanisms by which heat is transported within the solar plasma?

A

Heat is transported within the solar plasma through radiative losses, which include continuum and line emissions, and thermal conduction.

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22
Q

What is radiative instability, and how does it relate to the sudden temperature jump observed in the transition region of the solar atmosphere?

A

Radiative instability occurs when the rate of energy input into the plasma does not balance with the rate of energy output through cooling. In the transition region, a sudden temperature jump occurs due to the instability caused by the balance between heating and cooling mechanisms.

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23
Q

What assumptions are made in radiative instability

A

We assume the plasma is optically thin and that the cooling is dominated by photon emission.

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24
Q

How is the stability of the plasma system determined in the radiative instability theorem?

A

The stability of the plasma system is determined by the sign of derivative of the radiative loss function, f’. If f’ > 0 it is stable and f’ < 0 it is unstable.

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25
Q

What is the significance of the equilibrium temperature T1 in the transition region model?

A

The equilibrium temperature T1 marks the point where the equilibrium temperature abruptly transitions to the next accessible equilibrium temperature, T2, as density decreases and radiative losses increase higher in the atmosphere.

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26
Q

What role does thermal conduction play in defining the transition region of the solar atmosphere?

A

Thermal conduction takes over from radiation in defining the transition region, leading to a steep jump in temperature. It transports energy backward towards the lower chromosphere, contributing to the observed thickness of the transition region.

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27
Q

How is the heat flux conducted and how is it balanced?

A

The heat flux is conducted out of the transition region and balanced by a flux of power into the transition region. The exact nature of which is still under investigation.

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28
Q

What is the significance of downward heat conduction in the corona?

A

X-ray observations indicate downward heat conduction in the corona which contributes to the observed steady radial increase in temperature.

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29
Q

What are two observational constraints that must be satisfied by any solar corona model?

A

Observations indicate that in the outer layers of the solar atmosphere 1) density decreases
2) temperature is approximately constant

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30
Q

Can a static equilibrium be achieved in the solar corona?

A

No, a static equilibrium cannot be achieved in the solar corona.

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31
Q

What is the problem with the isothermal static corona model?

A

The isothermal static corona model fails because as the distance from the Sun increases, the density approaches zero, which is not physically plausible and contradicts observations.

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32
Q

How does the non-isothermal static atmosphere model attempt to address the issue?

A

The non-isothermal static atmosphere model considers temperature variations with radius to establish a more realistic profile. However, this model also fails to provide a satisfactory solution.

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33
Q

What is the Chapman model, and how does it relate to the solar corona?

A

The Chapman model suggests that the outer corona is heated by thermal conduction from the inner corona. However this model does not work.

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34
Q

Why is pressure from the interstellar medium (ISM) unable to contain the solar corona?

A

Pressure from the interstellar medium is insufficient to contain the solar corona due to the extreme difference in pressure between the Chapman model and the ISM.

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35
Q

What conclusion did Parker reach regarding the solar corona’s equilibrium and what did he identify?

A

Parker concluded that the solar corona cannot be in hydrostatic equilibrium and identified the coronal expansion flow gives rise to the solar wind.

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36
Q

What are some characteristics of the solar wind according to the Parker model?

A

According to the Parker model, the solar wind exhibits characteristics of a steady, spherically symmetric mass flux, with velocity profiles known as “isothermal solar winds.”

37
Q

What is the significance of the critical radius in the context of the solar wind?

A

The critical radius is the point at which the wind speed equals the sound speed, and beyond this point, the velocity profile of the solar wind changes.

38
Q

What are the different types of solutions based on the value of the constant in the context of solar wind models?

A

There are six sets of solutions categorized into different types. Types 5 and 6 are not continuous in velocity and are ignored. Types 2 and 4 start supersonic and are disregarded as they do not match observations. Type 3 starts and ends subsonic and is also ignored. Type 1 starts subsonic and ends supersonic, matching observations, and is considered critical.

39
Q

What is the significance of Type 1 solutions in solar wind modeling?

A

Type 1 solutions, which start subsonic and end supersonic, are significant because they match observations and pass through a critical point where the wind becomes supersonic.

40
Q

How is the solar wind density profile determined in the Parker solution, and what does it show? Additionally, what is this consistent with?

A

The solar wind density profile is determined from the mass continuity equation, showing that density decreases to zero at large distances. This density profile is consistent with observations and a low interstellar medium (ISM) pressure.

41
Q

What is a stellar breeze, and how does it relate to solar wind modeling?

A

A stellar breeze is related to a specific solution (set 3). In this scenario, mass continuity allows for a constant non-zero density and pressure. Such that a stellar breeze solution cannot be confined by the interstellar medium.

42
Q

What conditions lead to the solar wind becoming supersonic?

A

The solar wind becomes supersonic where plasma pressure starts to dominate over gravity, typically beyond a certain critical radius.

43
Q

What assumptions are made in the discussion of the Parker solution for solar wind modeling?

A

In the Parker solution, assumptions include considering an isothermal wind, disregarding other mechanisms driving the solar wind (e.g., radiation pressure), and ignoring the effects of the Sun’s magnetic field.

44
Q

How is the solar wind speed categorized, and what are the characteristics of each category?

A

The solar wind speed is categorized into slow and fast solar wind. The slow wind speed aligns with expectations from the Parker solar wind model, with a plasma temperature around 10,000 Kelvin. The fast solar wind, however, exceeds the speeds predicted by the Parker model.

45
Q

What is ram pressure, and why is it significant in solar wind dynamics?

A

Ram pressure is the dynamic pressure due to bulk motion. It affects the behaviour of the wind when it encounters obstacles or other plasma with different velocities and pressures.

46
Q

Why is the dynamic pressure important?

A

The dynamic pressure holds back the ISM.

47
Q

What effect does the mass loss due to pressure-driven winds have on stellar evolution?

A

the mass loss due to pressure-driven winds is too small to have any effect on stellar evolution.

48
Q

How is the mass-loss rate for pressure-driven winds determined?

A

The mass loss rate for pressure driven winds is given when the wind becomes supersonic (when the coronal density is at the critical point).

49
Q

When do high mass loss rates occur in pressure-driven winds?

A

High mass loss rates occur when the effective gravity acting on the wind is less, this is due to an outward radiation pressure.

50
Q

How does radiation pressure affect the mass-loss rate in stellar winds and where is this most pronounced?

A

Radiation pressure can significantly increase the mass-loss rate by “loosening” the grip of gravity on wind particles. This effect is particularly pronounced in early-type stars with high luminosities.

51
Q

What assumptions are made for radiation-driven winds?

A

Radiation-driven winds assume that the wind is composed of particles with a constant mass and cross-section, and that the mass flow is radial. It also assumes that the wind is optically thin, allowing all particles to be affected by radiation from the star.

52
Q

What is the outward radiation force?

A

The outward radiation force is an inverse-square law with an opposite action to that of gravity.

53
Q

What is the significance of the Eddington luminosity in the context of radiatively-driven winds?

A

The Eddington luminosity represents the point at which a star becomes radiatively unstable. This threshold plays a crucial role in determining the conditions under which radiatively-driven winds occur.

54
Q

What complications arise in modeling radiatively-driven winds?

A

Several complications arise, including the inability for wind velocity to be zero at the stellar surface unless the density there is zero, the non-point source nature of the star affecting radiation pressure, and the non-grey nature of the atmosphere, with scattering cross-sections dependent on wavelength.

55
Q

What are P-Cygni profiles, and why are they significant in stellar wind diagnostics?

A

P-Cygni profiles are characteristic line profiles observed in line-driven winds with strong resonant lines, named after the P-Cygni star. These profiles are particularly important when there is significant resonant scattering.

56
Q

How is the shape of P-Cygni profiles influenced by scattering processes?

A

The shape of P-Cygni profiles is influenced by two main scattering processes: Doppler broadening due to the expansion of the wind shell and scattering of photons from other regions into the line of sight, resulting in both absorption and emission features within the line profile.

57
Q

What is the significance of inverse P-Cygni profiles, and in what type of stars are they observed?

A

Inverse P-Cygni profiles can be observed in stars with substantial accretion discs, such as T-Tauri stars. These profiles are significant as they indicate a reversal in the typical absorption and emission features seen in P-Cygni profiles and can switch between the two types over short time periods.

58
Q

Why must we use a multiple scattering argument?

A

Single scattering often leads to an underestimation of the terminal velocity.

59
Q

What are the different types of stellar winds discussed?

A
  1. Coronal (pressure-driven) winds
  2. Dust-driven (radiation-driven) winds
  3. Line-driven (radiation-driven) winds
  4. Pulsation-driven winds
  5. Magnetically-driven winds
  6. Magnetically-rotating winds
60
Q

What are coronal winds and name an example?

A

Coronal winds are driven by gas pressure from a hot corona. The solar wind is a coronal wind.

61
Q

What are dust-driven winds?

A

Dust-driven winds are when radiation supplies the force to the dust. Dust in the outer envelopes drags the plasma out with it.

62
Q

What are line-driven winds and what is it effected by?

A

Line-driven winds is when radiation pressure is exerted via absorption and scattering of photons in strong resonant spectral lines. It is effected by the velocity profile and doppler shift.

63
Q

What are pulsation-driven winds?

A

Pulsation-driven winds are found in cool variable stars, where radial pulsations lead to the ejection of outer layers.

64
Q

What are magnetically-driven winds?

A

Magnetically-driven winds is when plasma flows along magnetic field lines, resulting in Alfven waves traveling along the field lines.

65
Q

What are magnetically-rotating winds?

A

Magnetically-rotating winds is when plasma flows along magnetic field lines and is flung out by the rotation of the star, resulting in Alfven waves travelling along the field lines.

66
Q

What is the Maxwell-Faraday law?

A

It states that a circulating electric field is produced by a magnetic field that changes with time.

67
Q

What is Lenz’s law?

A

Lenz’s law states that currents induced by changing magnetic flux always flow in a direction so as to oppose the change in flux.

68
Q

What is flux freezing in the context of heliophysics?

A

Flux freezing known as Alfven’s theorem, states that in a fluid with infinite electrical conductivity, the magnetic field is frozen into the fluid and moves along with it. The magnetic field and plasma are frozen together.

69
Q

What does Ampere-Maxwell’s law tell us about currents in magnetic flux tubes?

A

Ampere-Maxwell’s law states that currents traveling in opposite directions produce a repulsive force, while currents traveling in the same direction produce an attractive force.

70
Q

What is the significance of the plasma beta parameter in heliophysics?

A

Plasma beta (β) is a dimensionless parameter that determines the dominance between magnetic field and plasma motions. β &laquo_space;1 implies the magnetic field dictates plasma motion, while β&raquo_space; 1 implies plasma dominates and distorts the magnetic field.

71
Q

How does the plasma behave in the solar corona in terms of plasma beta?

A

In the solar corona, β &laquo_space;1, indicating that the plasma streams along magnetic field lines.

72
Q

Describe the behavior of the plasma near the Earth in terms of plasma beta.

A

Near the Earth, the magnetic field acts as a tracer of where the plasma is, indicating β ≳ 1.

73
Q

What are the two regimes of plasma beta in relation to heliophysics?

A

β &laquo_space;1 applies to the solar corona where the magnetic field dominates, while β&raquo_space; 1 applies to the solar wind far from the Sun, where plasma dominates and distorts the magnetic field.

74
Q

What are the three mechanisms through which the Sun influences the environment in space?

A

1) Fast solar wind from coronal holes.
2) Slow solar wind from edges of active regions and coronal streamers.
3) High-speed transients associated with solar flares and coronal mass ejections.

75
Q

How does space weather occur?

A

Space weather occurs when magnetic fields, plasma, and energetic particles ejected from the Sun interact with the Earth’s magnetosphere and atmosphere, producing various effects.

76
Q

What causes the variability of the solar atmosphere?

A

The variability of the solar atmosphere is caused by changes in magnetic fields at different time scales, including the 11-year solar cycle, solar rotation, active regions, and transient phenomena like flares and coronal mass ejections.

77
Q

What is magnetic reconnection and how is it related to Alfven’s theorem?

A

Magnetic reconnection is a process believed to power solar flares and coronal mass ejections, where magnetic topology is rearranged, and magnetic energy is converted to kinetic energy, thermal energy, and particle acceleration. It is a break down of Alfven’s theorem.

78
Q

Describe the characteristics of open and closed magnetic field lines in the solar atmosphere. Additionally, describe current sheets.

A

Open field lines extend to large distances and originate from regions of lower plasma density where you find coronal holes.

Closed field lines confine plasma to higher densities.

Current sheets form above the cusp of coronal streamers.

79
Q

Describe the solar wind in the rotating coordinate system.

A

In the rotating coordinate system, the solar wind velocity components include a radial component and an azimuthal component due to the transformation to the rotating frame of reference.

80
Q

What is the significance of the Alfven surface in solar wind dynamics?

A

The Alfven surface marks the boundary where the solar wind speed becomes super-Alfvenic, meaning it exceeds the Alfven speed. This boundary divides regions where the plasma follows the co-rotating magnetic field and where the magnetic field is dragged by the plasma forming Parker spiral.

81
Q

What are the main types of fluctuations observed in the solar wind and where are they greatest?

A

Fluctuations in solar wind velocity associated with waves and turbulence are observed throughout much of the solar wind. These fluctuations tend to be greatest in high-speed streams and are often Alfvénic in nature, meaning they involve coupled changes in flow velocity and magnetic field vectors.

82
Q

How do Alfven waves dissipate their energy in the solar wind?

A

Alfven waves propagate away from the Sun and dissipate their energy via a turbulent cascade, leading to heating and wind acceleration.

83
Q

What is the energy spectrum of turbulence?

A

The energy spectrum of turbulence, 𝐸(𝑘), is related to the mean turbulence kinetic energy per unit mass. It represents the contribution to turbulence kinetic energy by wavenumbers from 𝑘 to 𝑘 + 𝑑𝑘.

84
Q

What is the Kolmogorov spectrum and how does it relate to turbulence energy transfer rate?

A

The Kolmogorov spectrum is a universal spectrum for turbulence, independent of large-scale driving and small-scale dissipation. The cascade or energy transfer time in the Kolmogorov spectrum is proportional to 𝑘^(-2/3), where 𝑘 is the wavenumber.

85
Q

What is the Iroshnikov-Kraichnan spectrum and how does it differ from the Kolmogorov spectrum?

A

The Iroshnikov-Kraichnan spectrum is flatter than the Kolmogorov spectrum. It suggests a longer interaction time compared to the Kolmogorov spectrum and implies an anisotropic MHD (magnetic hydro dynamics) cascade.

86
Q

How does the Goldreich-Sridhar theory explain turbulence in the solar wind?

A

The Goldreich-Sridhar theory suggests that there is a critical balance between perpendicular and parallel scales in MHD turbulence, leading to anisotropic turbulence with scale-dependent anisotropy.

87
Q

What role does MHD turbulence play in solar phenomena?

A

MHD turbulence is believed to be crucial for explaining phenomena such as coronal and solar wind heating, acceleration of the solar wind, solar flares, coronal mass ejections, and magnetic reconnection at the Sun.

88
Q

Why do line driven winds depend on the Doppler effect?

A

If it weren’t for the Doppler effect the outer layers of hot stellar atmospheres would not be able to absorb any photons coming from the star.

89
Q

What is scale height?

A

Scale height of an atmospheric layer expresses how quickly density will drop

Heliophysics and Stellar Atmospheres (4 decks)

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