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Meteorology ENV 2A23

Meteorology ENV 2A23. Radiation Lectures. How is energy transferred?. How is energy transferred?. Conduction Convection Radiation. Conduction Convection Radiation. Conduction Convection Radiation. How is energy transferred?. Conduction – energy transfer from molecule to molecule

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Meteorology ENV 2A23

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  1. Meteorology ENV 2A23 Radiation Lectures

  2. How is energy transferred?

  3. How is energy transferred?

  4. Conduction • Convection • Radiation

  5. Conduction • Convection • Radiation

  6. Conduction • Convection • Radiation

  7. How is energy transferred? • Conduction – energy transfer from molecule to molecule • Convection – spatial mixing of “air parcels” i.e. masses of air • Radiation – primary source of energy for the Earth Radiation imbalances drive the circulation of the atmosphere and ocean

  8. Electomagnetic radiation in the range 0.1 to 10 micrometres (mm), i.e. 0.1-10 x10-6 m

  9. Electomagnetic radiation travels in packets (quanta), whose energy is given by E = hc/8, where 8 is wavelength, h is Planck’s constant (6.625x10-34 J s-1) c is speed of light (3x108 m s-1)

  10. The Sun • Most solar radiation is emitted from the photosphere (T~6000 K) • Sun powered by nuclear fusion, H to He • Plasma ejected as “solar wind”

  11. The Sun • The sun’s radiative output is centred on visible wavelengths

  12. The Sun • The sun’s output is not constant • Sunspot cycle ~11 years • Periods of high/low activity

  13. Sun-Earth Geometry • Axial tilt = 23.5o • Eccentricty = 0.02 • Aphelion = 1.50x108 km, 3 July • Perihelion = 1.45x108 km, 3 January • SH receives more solar radiation in summer than NH • Is it warmer?

  14. Sun-Earth Geometry • Equinoxes = “equal” days and nights

  15. Sun-Earth Geometry • Solstice = “sun stands still”, longest/shortest days

  16. Changes in orbital parameters result in changes in incoming solar radiation and distribution (Milankovitch 1930)

  17. The Sun’s energy output • The solar constant is the radiation flux density at the top of the atmosphere, for the mean sun-earth distance • i.e. the amount of radiation falling on the top of the atmosphere (per unit area) • S0 = 1360 W m-2

  18. The Sun’s energy output • The sun is an almost perfect emitter of radiation, i.e. emits maximum possible radiation for its temperature • It is a blackbody emitter and so governed by Stephan-Boltzmann Law: F = FT4, where, F is flux density W m-2, T is temperature, F = 5.67x10-8 W m-2 K-4

  19. Radiation flux density at the Earth • F = FT4 per unit area • So over sphere 4Brs2FT4 • Hence at distance of earth (rd): 4Brs2FT4/ 4Brd2 • i.e. S0 = rs2/rd2FT4, an inverse square law sun rs rd earth

  20. Emission temperature of a planet The emission temperature of a planet is the blackbody temperature with which it needs to emit radiation in order to achieve energy balance. To calculate this for the Earth, equate blackbody emission with amount of solar energy absorbed. - see radiation practical

  21. Emission temperature of a planet Energy incident on planet= solar flux density x shadow area But not all radiation is absorbed, some is reflected: albedo (α) = reflected/incident radiation Absorbed solar radiation = S0(1- α)π re2 (W) Absorbed solar radiation per unit area = S0(1- α)/4 (W m-2) This must be balanced by terrestrial emission. If we approximated Fe as a blackbody: FEarth = σTe4 , where Te is the blackbody emission temperature. => Te4 = S0(1- α)/σ4 For Earth, Te = 255 K. Note this is well below the average surface air temperature of the Earth = 288 K.

  22. Distribution of Insolation • Seasonal & latitudinal variations in temperature are driven primarily by variations in insolation • The amount of solar radiation incident on the top of the atmosphere depends on:

  23. Distribution of Insolation • Seasonal & latitudinal variations in temperature are driven primarily by variations in insolation • The amount of solar radiation incident on the top of the atmosphere depends on: • Latitude • Season • Time of day

  24. Distribution of Insolation The solar zenith angle (2s) is the angle between the local normal to the Earth’s surface & the line between the Earth’s surface & the sun The (daily) solar flux per unit area can be calculated as: where S0 is the solar constant, and d is the sun-earth distance 2s earth

  25. Distribution of Insolation • The season ~ declination angle *, • i.e. latitude on Earth’s surface directly under the sun at noon - * varies between 23.5 & -23.5o • The time of day ~ hour angle h, • Longitude of subsolar point relative to its position at noon • Then cos θs = sinφ sinδ + cosφ cosδ cosh, for latitude φ

  26. Distribution of Insolation

  27. Distribution of Insolation • Equator receives more solar radiation than the poles (at the top of the atmosphere)

  28. Energy balance at the top of the atmosphere • As well as the distribution of insolation, the amount of energy absorbed and emitted depends on atmospheric and surface conditions.

  29. Energy balance at the top of the atmosphere • albedo (α) = reflected/incident radiation

  30. Energy balance at the top of the atmosphere • Outgoing longwave radiation

  31. Energy balance at the top of the atmosphere Net radiation The net radiation can be calculated from R = SWd – SWu + LWd – LWu , Where SW = shortwave (solar) radiation, LW = longwave (terrestrial radiation) => R = SWd(1-αp) –LWu at the top of the atmosphere, where αp is the planetary albedo.

  32. Energy balance at the top of the atmosphere => R = SWd(1-αp) –LWu at the top of the atmosphere, where αp is the planetary albedo.

  33. Energy balance at the top of the atmosphere There must be a poleward transport of energy to balance out the net gain at the equator and the net loss at the poles.

  34. Radiation Flux and Radiation Intensity • The radiation flux density (or irradiance), F (units W m-2) is the radiant energy crossing a unit area in unit time. It does not discriminate between different directions. • The radiation intensity (or radiance), I, (units W m-2 steradians-1) includes information on directionality. • Special Case : • Radiation intensity I is isotropic, • Then F = BI • For example: emission from a blackbody, emission from the atmosphere Animation…

  35. What about the wavelength of the radiation? Planck’s Law • In other words, radiation intensity depends on frequency (or equivalently wavelength) of emission. i.e. B Bv(T) dv = FT4

  36. What about the wavelength of the radiation? Wein’s Law Sun’s emission peaks ~ 4.8 microm Earth’s emission peaks ~ 10 microm Brightness temperatures of the sun and Earth are ~6000 K and 255 K

  37. What about the wavelength of the radiation? When an object is not a blackbody, then its radiation flux density can be written F = eσT4, where e is the emissivity. Usually eλ = e(λ) is a function of wavelength. If we define absorptivity aλ as the fraction of incident radiation that is absorbed. It can be shown that eλ = aλ , this is Kirchoff’s Law. i.e. an object emits radiation at each wavelength as efficiently as it absorbs it.

  38. Radiation in the atmosphere • Earlier we found the blackbody emission temperature Te = 255 K, much colder than the observed Tsurface = 288 K. • Why ?

  39. Radiation in the atmosphere • Difference is due to selectivescattering, absorption and emission of radiation by the atmosphere. • These depend upon the structure of the molecules present. sketch

  40. Radiation in the atmosphere • Difference is due to selectivescattering, absorption and emission of radiation by the atmosphere. • These depend upon the structure of the molecules present.

  41. Scattering • Scattering decreases the intensity of the solar beam. • It depends upon λ (wavelength) and d (particle size). • Three cases:

  42. (1) Rayleigh Scattering occurs when d << λ For example from O2 or N2, the major tropospheric gases, where d = 10-10 m and λ = 0.5x10-6 m. Scatters equal amounts of radiation forward and backward The amount of scattering strongly dependent on λ: the volume extinction coefficient is a function of 1/ λ4 Rayleigh scattering explains why the sky is blue and sunsets are red. - blue (short λ) scattered more than red (long λ) light

  43. (2) Diffuse scattering occurs when d >> λ • Diffuse scattering occurs when d >> λ, for example from dust or cloud droplets • Typically ~10 mm • Diffuse scattering is independent of λ. • Clouds appear white and polluted skies are pale • Full consideration requires Mie theory.

  44. (3) Complex Scattering occurs when d = λ • Diffraction

  45. Absorption • All gases absorb and re-radiate energy at specific wavelengths depending on their molecular structure. • Electronic excitation – visible uv • Vibrational excitation – IR • Rotational excitation – thermal IR • Molecules need a permanent electric dipole, e.g. H2O O - H H +

  46. Aborption occurs at specific wavelengths (lines) according to the excitational properties of the gas (or gases) involved. • However these lines are broadened by various mechanisms into absorption bands.

  47. Absorption line broadening • Natural broadening – associated with the finite time of photon emission and the uncertainty principle • Pressure broadening (or collision broadening) – collisions between molecules supply or remove small amounts of energy during radiative transitions. - Primary mechanism in the troposphere (why?) • Doppler broadening – results from the movement of molecules relative to photons. - dominant at higher altitudes

  48. Groups of lines within a frequency interval are termed absorption bands • In the thermal infra-red there are important absorption bands due to H2O, CO2, O3, CH4, N2O, etc

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