Infrared Solar Physics

Abstract

The infrared solar spectrum contains a wealth of physical data about our Sun, and is explored using modern detectors and technology with new ground-based solar telescopes. The scientific motivation behind exploring these wavelengths is presented, along with a brief look at the rich history of observations here. Several avenues of solar physics research exploiting and benefiting from observations at infrared wavelengths from roughly 1000 nm to 12 400 nm are discussed, and the instrument and detector technology driving this research is briefly summarized. Finally, goals for future work at infrared wavelengths are presented in conjunction with ground and space-based observations.

Keywords: Detectors; Solar atmosphere; Solar magnetic fields.

Figure 1:

Atmospheric transmission from 560 to 21000 nm as measured from the McMath-Pierce facility at Kitt Peak. The molecules which are responsible for the various absorption bands are marked. Image reproduced with permission from Hinkle et al. (2003), copyright by AAS.

Figure 2:

Atmospheric scattered light as a function of wavelength from sea level measurements over water, taken in 1959–1960. Corrections for water vapor and methane absorption were made. The wavelength behavior of the scattering changes on different days, presumably depending on the particulates present during the testing. These particular measurements show at best a λ^−1.7 change, but at other times are nearly flat. Image reproduced with permission from Knestrick et al. (1962), copyright by OSA.

Figure 3:

A schematic representation of the wavelength variation of the Mueller matrix of the DST. In each box the wavelength dependence of the value of the matrix element is plotted from 400–1600 nm. What is important to note is that the off-diagonal elements, which represent the polarization cross-talk, approach a value of 0 as the wavelength increases. Image reproduced with permission from Socas-Navarro et al. (2011), copyright by ESO.

Figure 4:

The benefit of making magnetic field Zeeman observations in the infrared is shown in this figure. On the left is a plot showing the separation of the Stokes V peaks for various Voigt profiles at a magnetic field of 1 kG for a visible line and an IR line. On the right is a plot showing the value of the Stokes V amplitude. The Zeeman splitting of the 1565 nm IR line is fully resolved at this magnetic field strength, and the weak field approximation does not need to be used to interpret the spectra. Image reproduced with permission from Stenflo et al. (1987), copyright by ESO.

Figure 5:

The height of formation of the infrared continuum radiation is shown in this figure. The continuum radiation is formed at the deepest level in the Sun at 1600 nm, and then the continuum formation height increases with increasing wavelength in the infrared spectrum. The bottom plot shows that the dominant source of opacity is from Hydrogen free-free transitions. See also Fontenla et al. (2006) for more recent work. Image reproduced with permission from Vernazza et al. (1976), copyright by AAS.

Figure 6:

This figure shows the types of changes needed in the solar atmosphere models in order to bring the CO observations into agreement with other chromospheric measurements. The diagram on the left shows an early simple atmospheric model, and the diagram on the right shows the updated model after the discovery of strong CO lines and dynamic shock waves in the chromosphere. Image reproduced with permission from Ayres (2002), copyright by AAS.

Figure 7:

The limb extensions in arcseconds of several CO lines near 4666 nm are shown in this plot. Image reproduced with permission from Ayres and Rabin (1996), copyright by AAS.

Figure 8:

Diagnostic l-ν diagrams for full disk oscillations observed with the Doppler shifts and depths of CO 4666 nm lines. The CO oscillations show the same power ridges as the global p-mode oscillations, but show an I–V phase which is different from the expected adiabatic phase shift and which allows an investigation of the dynamics of the solar atmosphere. Image reproduced with permission from Penn et al. (2011), copyright by AAS.

Figure 9:

A figure showing the polarization spectra of a sunspot. The spectrograph slit crossed a sunspot umbra, and the Zeeman spitting in Stokes I, Q, U, and V are shown in this figure, with the continuum intensity removed from the Stokes I panel. The 1565 nm g = 3 line is shown at the center of each panel, and other atomic and molecular lines are identified in each spectrum. Image reproduced with permission from Penn et al. (2004a), copyright by AAS.

Figure 10:

Fits to the Stokes spectra for the 1565 nm Fe line pair. Here the Fe and OH lines are both fit using response functions computed for each line. Image reproduced with permission from Mathew et al. (2003), copyright by ESO.

Figure 11:

A full disk spectroheliogram in He I 1083 nm from the NSO/VSM instrument; solar north is up. The image shows polar coronal holes as regions of less absorption, and a low latitude coronal hole in the northern hemisphere. Dark absorption accompanies active regions, and the quiet Sun shows less absorption, but reveals the internetwork pattern. Limb prominences and limb emission are seen as bright regions off the solar limb, and a dark filament is visible on the disk in the northern hemisphere. Credits: NSO/AURA/NSF.

Figure 12:

A plot of flare emission seen in the He I 1083 nm line. The bottom spectrum shows very strong absorption in an active region filament, the central spectrum shows strong active region absorption, and the top spectrum shows emission in a flare kernel. In all cases both He I line components are seen. Image reproduced with permission from Penn and Kuhn (1995), copyright by AAS.

Figure 13:

A composite spectrum of a strong He I 1083 nm downflow. At every pixel, the normal spectral line at the rest wavelength is seen along with a highly red-shifted second spectral line. The Stokes IQUV spectra are taken from along a curved path following a coronal condensation event. The upper panels show the measurements, while the lower panels show the resulting fitted spectra from the analysis, with the telluric and weak photospheric lines removed. The magnetic field measured in this event varies with height from approximately 100 G at 50 Mm to over 1000 G at 10 Mm and below. Image reproduced with permission from Schad (2013), copyright by the author.

Figure 14:

Oscillations observations in He I 1083 nm. Both velocity and intensity oscillations were seen, but power peaks are only visible in the velocity signal. The coherence is high at low frequencies, but the phase shift is very different from the adiabatic value of 90 degrees. Image reproduced with permission from Fleck et al. (1994), copyright by IAU.

Figure 15:

Simultaneous spectra from a sunspot umbra showing the Mg I 12318 nm (top) and the Fe I 1565 nm (bottom) lines. The larger splitting of the Mg I line components is clearly visible, and the width of the Zeeman shifted components reveals a large range in the magnetic field strengths seen at this spatial position. Image reproduced with permission from Moran et al. (2000), copyright by AAS.

Figure 16:

Maps of the magnetic field strength in a magnetically complex active region. The Zeeman shifted sigma components of the line are sliced at different wavelength positions corresponding to the magnetic field values listed in each sub-image. Magnetic fields of various strengths are seen to occur in the same spatial pixels, while the stronger magnetic fields are limited to the umbral regions. Image reproduced with permission from Jennings et al. (2002), copyright by AAS.

Figure 17:

A map of the ratio of emission from the two infrared [Fe XIII] lines near 1075 nm from the 1991 eclipse. The line ratio is sensitive to the local coronal electron density and reveals structures in the hot plasma near a cool prominence, seen in the lower right of this figure. Image reproduced with permission from Penn et al. (1994), copyright by AAS.

Figure 18:

The coronal spectrum from 1000 to 1500 nm of the 1994 eclipse. The left panel shows the raw counts, and the right panel is normalized to the disk center solar intensity. Lines are seen at 1075, 1080, 1083, 1251, and 1431 nm. Image reproduced with permission from Kuhn et al. (1996), copyright by AAS.

Figure 19:

A spectrum of the 3934 nm [Si IX] coronal emission. The top panel shows the observed coronal spectrum with the emission line fit with a Gaussian line profile; the middle panel shows the coronal spectrum and a disk spectrum for reference; and the bottom panel shows a graph approximating the expected Stokes V profile for this line. Image reproduced with permission from Judge et al. (2002), copyright by AAS.

Figure 20:

The coronal magnetic field variation with height in the corona over an active region at the solar limb. Measurements are shown (with error bars) as the solid line, and the polarity of the observed magnetic field is seen to vary. Expected magnetic field values from an extrapolation of the active region fields are shown with the points. Image reproduced with permission from Lin et al. (2004), copyright by AAS.

Figure 21:

A map of the outgoing and ingoing power seen in 5-minute period coronal waves for a section of the corona. The oscillations are measured with the [Fe XIII] 1075 nm Doppler shift, and the waves are classified as outgoing and ingoing using their propagation direction. The outgoing waves are measured to have more power than the ingoing waves. Image reproduced with permission from Tomczyk and McIntosh (2009), copyright by AAS.

Figure 22:

Images of solar granulation in the infrared J, H, K bands. In each case a sequence of short exposure images was processed using the Multi-Object Multi-Frame Blind Deconvolution (MOMFBD) reconstruction algorithm. In some images artifacts from the reconstruction are visible. The images are not from the same regions of the Sun and are taken at different times under different seeing conditions. The axes are marked in units of arcseconds. Comparison of the contrasts at these wavelengths shows larger than expected contrast in the K-band images. Image reproduced with permission from Penn (2008), copyright by AGU.

Figure 23:

The Zeeman Stokes V splitting patterns for several molecular lines. Although the molecular lines measure the same magnetic field, the properties of the particular lines produce positive and negative polarity Stokes V profiles. Image reproduced with permission from Harvey (1985), copyright by NASA.

Figure 24:

The Stokes profiles for the 1526 nm Mn I absorption line for various magnetic field strengths are shown in this figure. For magnetic fields less than about 1000 G, the two spectral components in the Mn I line show different strengths, and then for fields at higher values the spectral components more apart in wavelength. The special hyperfine structure involved in this line provides a uniquely sensitive measurement of the solar magnetic field. Image reproduced with permission from Asensio Ramos et al. (2007), copyright by AAS.

Figure 25:

Stokes V spectral profiles of a Ti I line at 2231 nm. The two Zeeman split components are completely resolved in the sunspot umbra, and the line does not form in the high temperatures of the solar photosphere. With a magnetic sensitivity larger than the Fe I 1565 nm spectral line, this line is uniquely suited for probing the magnetic fields in sunspot umbrae. Image reproduced with permission from Penn et al. (2003b), copyright by Springer.

Figure 26:

Stokes I and V profiles for several atomic and molecular lines near 4135 nm from NAC data taken at the McM-P telescope. Credits: NSO/AURA/NSF.