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• Outline
• Glossary
• 7. Observing Systems
  • Previous: 7.6 Polarization and Antenna Feeds
  • Next: 7.8 Radio Frequency Interference (RFI)

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7.7 Propagation effects

Cosmic radio waves traverse a variety of ionized and neutral media before measurement at a radio telescope. Some effects are negliable and/or easily corrected for, however, others can prevent any useful astrophysical inference at all regardless of the sensitivity of the telescope. This section will provide a brief overview of the dominant sources of signal corruption in the troposphere, ionosphere and interstellar media.

7.7.1 Interstellar medium

The interstellar medium is filled with plasma that is not uniformly distributed, which causes radio waves to scatter and astrophysical sources to scintillate (vary in apparent brightness), increasingly so at low frequency. Here we cover some basic scattering theory, which is of course applicable to the other astrophysical media discussed here.

An electro-magnetic wave passing through any medium with refractive index inhomogeneity will cause the wave to scatter. This effect is most commonly modeled as a thin screen (relative the to the distance to the source) normal to the line-of-sight, which adds a random phase offset to the incoming wave at each transverse point, yielding what can be viewed as a 'corrugated' (not flat) wavefront.

Figure 7.7.1: Schematic diagram showing on a plane by a irregular slab of plasma. As can be imagined, what should be measured as constant phase across the radio telescope array, will now include a stochastic phase corruption due to differeing refractive indices (Image credit: Sanjay Bhatnagar PhD Thesis)

The degree of corrugation can be described by the phase structure function,

$$D_\phi (\mathbf{x},\mathbf{x'}) =\ < ( \phi(\mathbf{x} +\mathbf{x'}) - \phi(\mathbf{x}))^2>,$$

where $\mathbf{x}$ and $\mathbf{x'} $ represent two points on the screen. Assuming isotropic (no directional dependence), power law turbulence simplifies the expression to

$$D_{\phi} (r) = (r/r_{\rm diff})^{\beta},\ \ (r^2 = \mathbf{x}^2 -\mathbf{x'}^2),$$

where $r_{\rm diff}$, is the distance on the phase screen where the mean phase difference is equal to 1 radian. The exponent of $\beta = 5/3$ is characteristic of Kolmogorov turbulence, which empirical results have shown to be a reasonable approximation for the ISM and troposphere in many circumstances (see Thompson et al. 2001, and refernce therein).

Another important length scale is the Fresnel scale, defined as the distance across the phase screen where the geometrical path difference is $1/2$ radian, $r_{\rm F} = \sqrt{\lambda D/2\pi}$, where $D$ is the observer - scattering screen distance. These two length scales, $r_{\rm F}$ and $r_{\rm diff}$, define the nature of the scattering which is split into the \emph{strong} and \emph{weak} regimes. In weak scattering, $ r_{\rm diff} \gg r_{\rm F}$ and hence, $D_{\phi}(r_{\rm F}) \ll 1$. In the regime of \emph{strong scattering}, $ r_{\rm diff} \ll r_{\rm F}$ implying $D_{\phi}(r_{\rm F}) \gg 1$.

Using a frozen screen assumption, where micro-scale diffusion within the scattering screen is much smaller than the bulk motion, the structure function becomes $D(R)=D(t)|_{R=v T}$ where $v$ is the scattering screen transverse velocity and T is the integration time. This simply means one assumes the screen that causes the phase corruptions changes structure on time scales much larger than it takes for the for the screen to move across the telescope's field of view. This assumptions allows the structure function to be parameterise in terms of bulk velocity of the screen, rather than position. This has a number of advantages, as will become clear.

These length and timescales determine a wide range of scattering effects in radio imaging, which are summarized in Figure 7.7.2. A full treatment of these cannot possibly be covered here, however, Table 7.7.1 provides indicative values for a range of media, which can then be plugged into the ScatterBrane software below to test different regimes and gain a more intuitive understanding. The additional reading section at the end of the chapter will provide further references for more in-depth explanations.

Table 7.7.1: Typical values of scattering media in astronomy (Credit: Narayan 1992)

Figure 7.7.2: Different imaging regimes in the presence of scattering, shown as function of the integration time and Fresnel time ($t_{\rm F} = r_{\rm F}/v$) on the vertical axis and the strength of the scattering shown on the horizontal axis. These regimes can be explored using the ScatterBrane code below in combination with typical values from Table 7.7.1. (Credit: Narayan 1992)

The following code is taken from the SLIMSCAT package courtesy of Michael Johnson and Katherine Rosenfeld. Changing the input parameter values and input image will enable a more intuitive understanding of scattering effects in the different imaging regimes shown in Figure 7.7.2.

Figure 7.7.3: Example of the effect of strong scattering on a Gaussian source.

Clearly this degrades the quality of the observations significantly. Trying playing around with all the (hidden) input parameters to the slimscat.generate_screen and slimscat.run_slimscat functions. For more information see http://krosenfeld.github.io/slimscat. We finish with one of their fun examples, which is scattering Gargantua from the movie Interstellar. While this is simply illustrative, it is an real observation challenge that is faced by researchers trying to image the shadow of the black hole at the centre of our Galaxy. This is done for three colour filters.

Following the scattering, we need to smooth the image a little to better illustrate the result:

Make the RGB cube and you're done:

Figure 7.7.4: Fun colour example of a scattered RGB image of the black hole sillhouette of Gargantua (from the movie Interstellar).

If you increase the wavelength of the observation, or decrease the frequency (see input to slimscat.generate_screen), you'll see how interstellar scattering can be a major barrier to successfully measuring the shadow of a black hole with the Event Horizon Telescope! See http://www.eventhorizontelescope.org/ for more details on this project.

7.7.2 Ionosphere

The ionosphere is a layer of ionized plasma with an altitude range of 60-1000 km above sea level. This is driven mostly by ultraviolet photons from the Sun that ionize the upper layers of the atmosphere. As a result, there is a natural diurnal cycle, with the layers nearest the Sun being the mostly highly ionized. This level of ionization is measured with the unit Total Electron Content (TEC), where 1 TEC unit (TECU) = $10^{16}$ m$^{-2}$. Figure~xx shows a global map with the TEC values at a given time, showing that this is most concentrated near the equator.

Figure 7.7.5: Total Electron Content above the Earth at given time, showing the enhancement towards the equatorial regions nearest the Sun. (Image credit: http://www.astrosurf.com/luxorion/qsl-review-propagation-software-research.htm)

The ionosphere is a subject of intensive and complex study itself. Here we just describe the primary (adverse) effects it has on radio observations. Like any plasma, the ionosphere has a oscillation or plamsa frequency, given by $$\nu_p = \frac{e}{2\pi}\sqrt{\frac{n_e}{\epsilon_0 m}} \ \simeq 9\sqrt{n_e} \ \ {\rm (Hz)},$$ where $e, n_e, \epsilon_0$ are the electron charge, density and permitivity of free space respectively. A typical value for the ionosphere's plasma frequency is less than $\nu_p \sim 15$ MHz. Note than the ionosphere is opaque to any radiation below this frequency. While this disadvantage in low frequency observations of the Universe from Earth, it has also be used to aid long-distance radio communication by reflecting radio waves off the lower surface of the ionosphere.

The free electrons in the ionosphere cause a propogation delay of an incident electro-magnetic wave, which can be represented as an "excess path length", $$L_0 \simeq -\frac{1}{2} \int^\infty_0 \left[ \frac{\nu_p(h)}{\nu}\right]^2 {\rm d}h, $$

where $h$ is the height above sea level. Note the $\nu^2$ dependance, which is why the ionophere is primarily a problem for low radio frequency observations, such as attempts to detect the epoch of reionization. Within the Jones matrix formalism, this delay can be represented as follows,

$$ \mathbf{Z} = Z = \left[ \begin{array}{cc}\mathrm{e}^{-2\pi \imath \tau / \lambda} & 0 \\ 0 & \mathrm{e}^{-2\pi \imath \tau / \lambda} \end{array} \right] = \mathrm{e}^{-2\pi \imath \tau / \lambda}, $$

which is of course the same as any other phase only effect, of which there are many described in this chapter. Although the $\mathbf{Z}$ Jones matrix has no amplitude effect, it is important to note that the variable or differential delay imposed by the ionsphere result in a loss of coherence if data are averaged on timescales larger than the local dynamical timescale for the ionosphere (seconds to tens of minutes). It this is loss of coherence that also results in an amplitude loss.

Figure 7.7.5: Comparison of atmospoheric and ionospheric temperature and density profiles as a function of height above sea level.

For a given observing frequency, the extent to which the ionosphere corrupts the measured visibilies and the ease with which one can correct for that depends on the array geometry and station field of view (FoV). These can be best summarized as follows:

1. small array, small FoV: strong phase coherence across array and FoV
2. large array, small FoV: some/no phase coherence across array, strong phase coherence across FoV
3. large array, small FoV: strong phase coherence across array, some phase coherence across FoV
4. large array, large FoV: some/no phase coherence across array, some phase coherence across FoV

Actual scales for the above descriptions depend on the observing frequencies and ionospheric turbulence, however, at a few hundred MHz indicative values for "small arrays" can be for baselines less than 1 km; and an indicative value of a "small Fov" is < 10 arcmin$^2$. Cases 3 and 4 above demand that direction dependent calibration be employed to accurate solve for the ionospheric effects. This topic is under active research due to the slew of low frequency, wide FoV instruments under contruction or in their design phase.

Figure 7.7.6: Four different radio array regimes with regard to ionospheric calibration (see text). Credit: Huib Intema (2013 Lorentz Center presentation).

7.7.3 Trophosphere

This is the lowest layer of the neutral atmosphere and has typical scale height of 2 km. The refractivity of water vapour in this layer at radio wavelengths is of order 20x greater than in the optical/near-infrared (recall that typical "seeing" in the optical window is ~1 arcsec). Since water vapour is poorely mixed in this layer, there are phase fluctuations introduced due to the refractive index inhomogenities that result. The delay due to a non-unity refractive index ($n$) is $$\Delta t = \frac{1}{c} \int (n - 1) dy, $$ where c is the speed of light the refractive path length is the straight line path. The refractivity of moist air is described by the empirical formula $$N = 77.6\frac{p_D}{T} + 64.8\frac{p_V}{T} + 3.776 \times 10^5 \frac{p_V}{T^2}, $$ where the refractivity $N \equiv 10^6(n-1)$, $T$ is the temperature in Kelvins, $p_D$ and $p_V$ are the partial pressures of dry air and water vapor respectively. Since the atmosphere obeys the equation of hydrostatic equilibrium (to a good approximation), it is clear that to zero-th order, the trosphere's refractive index will have some dependence on time-of-day, time-of-season, site altitude and latitude. However, water vapour is not well mixed in the atmosphere and is therefore not well correlated with ground-based meteorological parameters.

Figure 7.7.7: Map of the Earth showing water vapour content during August 2010. (NASA image by Robert Simmon)

The troposphere has several effects on high frequency ($>> 10$~GHz) radio observations, the most important (or detrimental) are:

1. absorption
2. refraction
3. turbulence-induced scattering

The absorption (also known as the opacity) is due to pressure-broadened molecular resonances of water vapour, oxygen and ozone. Between 100 GHz to 1 THz there are a series of strong water vapour lines, necessitating the contruction of millimetre observatories at the high altitudes ($>2$ km), where several atmospheric windows are made accesible (see Figure 7.8.2).

The mean refraction is an effect that must be corrected in the telescope pointing, increasingly so at low elevation. However, small-scale tropospheric turbulence results in rapid phase fluctuations, significantly limiting telescope sensitivity. At $>$100 GHz, coherence times are limited to $\sim$10 seconds for even some of the best sites in the world.

The power spectrum of these phase fluctuations are usually described as a Kologmov turbulent screen, with a phase structure function defined as:

$$D_\phi(d) = \ <[\Phi(x) - \Phi(x - d)]^2>, $$

where $\Phi(x)$ is the phase at point $x$, $\Phi(x-d)$ is the phase at point $x-d$, and the angle brackets represent the ensemble average. This implies that the closer two points are to one another, the more similar their phase delay will be. The theory on this subject in particular is vast and not covered in any further detail here, but see Thompson et al. 2001, and reference therein.

In terms of the Jones matrix formalism, the troposhere is equiavalent to the ionosphere, however, with the additional amplitude loss due to non-zero atmospheric opacity.

Figure 7.7.8: Transmission curves (fraction of radiation that is able to traverse a medium) above the ALMA site Chajnantor. The lines correspond to the frequency coverage of the ALMA receiver bands (1-10). (Image credit: ESO http://www.eso.org/sci/facilities/alma/documents.html)

There are different regimes of atmosphereic phase corruptions, with the boundary conditions being two antennas looking through the exact same atmosphere, to VLBI baselines where the weather above two antennas would be entirely un-correlated. The below figure demonstrates the level of correlation in phase above the Very Large Array as a function of baseline length. The changes in slopes reveal the changing slope with the structure function.

Figure 7.7.9: Root phase structure function based on BnA array VLA 22 GHz observations. Open circles show the phase rms versus baseline length towards a 1 Jy calibrator source (0748+240) over the course of 90 minutes. The filled swquare show the same values with a constant noise value of 10$^\circ$ subtracted in quadrature, revealing the underlying tropospheric structure functions. All three regimes are predicted by Kolomogorov turbulence theory. (Image credit: Carilli & Holdaway 1999)

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• Next: 7.8 Radio Frequency Interference (RFI)