Instrumental optics and problems of the earth's atmosphere

Planet.

Space

Sci.,

1962, Vol.

9, pp. 675 to 699.

Pergamon

Press

Ltd.

Printed

in Northern

Ireland

INSTRUMENTAL OPTICS AND PROBLEMS EARTH’S ATMOSPHERE Astronomy

OF THE

A. H. JARRETT Section, Department of Applied Mathematics, The Queen’s University of Belfast (Received~2

July 1962)

Abstract-The suitability is discussed of spectrographic equipment for investigating various problems presented by the emissions in the upper atmosphere. Also an outline is given of instruments for synoptic and photometric studies of the airglow and aurora, and the merits of interferometric and spectrographic techniques set forth. Emphasis is on the practical aspect throughout, with the relevant historical background. INTRODUCTION

As recently as twenty five years ago it was quite usual for a spectrum of the nightglow to require tens or even hundreds of hours exposure time to achieve a satisfactory density. Even for the aurora, many of the more useful spectra were the result of an entire season’s observing. This arose simply from the fact that the prism type spectrographs were of poor lightgathering power and moreover the observer has to contend with photographic emulsions Nevertheless, many of the more intense radiations of the nightsky were of low sensitivity. identified with these relatively primitive instruments. However, problems of great theoretical interest, such as changes in the intensity of the emissions over short intervals of time, just could not be tackled. The spectra were merely the aggregate of all recorded variations for the duration of the exposure. A similar restriction applied to deductions of temperature from measurements of the spectra. Anyway, a fairly high resolving power is essential for accurate temperature derivations; the earlier spectrographs just did not possess this feature. As we shall see later, their dispersion was low, being of the order of several hundreds of A/mm in the red region of the spectrum. The development over the last decade of highly sensitive photoelectric photometers has opened up the possibility of synoptic studies aimed at acquiring a knowledge of the global Also the infra-red region has extent of the movements of aurora1 and nightglow features. been made accessible by progress in photoelectric detector techniques, and the ultra-violet by balloon, rocket, and satellite borne instruments. I. SPECTROGRAPHIC

(a) Slit-type

INVESTIGATIONS

OF THE NIGHTSKY

spectrographs

Considering the nightglow, we have of course the obvious difficulty, namely the feebleness of the radiation. As has been pointed out (l), the illumination level we have to deal with corresponds roughly to that of a candle held 100 m away. In the visual region of the spectrum we can assume that it generally falls below the visual threshold, although Roach and Jamnick@) report that nightglow structure can be seen at times. The infra-red radiation is much more intense, in fact if the eye responded to such emission the night sky would appear approximately as bright as it does during the time of full moon. In the case of the aurora the situation is rather better as the faintest activity is generally above the threshold value; for 675

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A. H. JARRETT

the 5577 A line the emission rate is around 1 kR for a weak aurora, and maybe 1000 kR for an intense displayt3). SlipheW was one of the first to record photograp~ca~y the 5577 line in the spectrum of the night sky. He used a small three prism spectrograph, with a camera lens of 15 in. focal length. Three years later, in 1922, Rayleighc5) published an account of his investigations of the green line in the spectrum of the night sky. The two instruments he used were of low dispersion, one giving 358 A/mm near the green line and the other about 240 ~~rnrn, the latter using a simple 60” prism as the dispersing element. The DalImeyer camera lenses worked at f/1*9,in conjunction with an achromatic collimator of 9 in. focal length. The plate scale of these slit-type spectrographs was quite sufficient for diagnosing the presence or otherwise of the green line. The slit width used was 0.12 mm. Too wide a slit could cause confusion in that the peak intensity in the continuous spectrum of the night sky might be mistaken for the aurora line itself. Rayleigh was also well aware of the fact that the low resolving power of his equipment could mean that the bright head of a band in reality composed what he recorded as line. It is interesting to note that in an attempt at visual observation of the 5577 line the slit was mounted in the shutter of a dark room and observed “through a powerful combination of direct vision prisms, which give tolerable spectroscopic purity with the slit wide. In this way, resting the eye in the dark for a time, the continuous spectrum of the night sky can be made out, and the slit may then be narrowed.” However, Rayleigh never succeeded in seeing the green line, and he attributed the perception of the continuum simply to the superior visibility in poor light of an object subtending a large angle at the eye. It was realized that high light-gathering power was essential for the success of these investigations. This implied short focal length camera lenses, with consequent low dispersion, if the prism dimensions were kept within practical limits. Rayleigh’s equipment required exposures of about four hours’ duration for a sufficiently dense impression of the green line. Aurora1 spectroscopy had its beginnings in 1867 when Angstrom@) found the 5577 line, along with several emissions in the blue and one in the red at 6300 A. The earlier work was carried out with direct vision instruments, but it proved very difficult to determine the wavelengths of the emissions accurately. Certainly the improvements in sensitivity of the photographic plate made things easier, but it is interesting to note that in the ~~~~~~e~ deu Spektroskopie of 1910 Kayser (‘) states that no reliable conclusion could be made as to their chemical origin. It was about this period that Vegard commenced his long series of researches on the spectrum of the aurora. For several years he used a sht-type spectrograph with an f/2*5 camera of 10 cm aperture, and a two-prism dispersing element. The dispersion was approximately 270 A/mm at 5000 A. The limitations imposed by the glass optics at the blue end of the spectrum occurred at 3400 A. To proceed down to the limit at 3100 A arising from atmospheric ozone cut-off it is essential to use quartz optics. Certainly an instrument of this type will record the stronger emissions (e.g. 5577, 6300, 6364 and 3914 A) within one hour from a class II (on international brightness coefficient scale(*)) display, but for study of the fainter features exposure times a~umulating from activity over several months are necessary. Exposing a spectrograph for long periods presents various problems in itself. Perhaps the most important is that the entire optical system must be adequately thermostatted. Otherwise definition will be adversely affected by changes induced by temperature in the

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dispersive properties of the prisms and the camera focal length. This latter point is of paramount importance, as we shall see later, in high definition Schmidt cameras. Fairly obvious, but of vital necessity for optimum results, is that the mechanical construction of the instrument must be sound. No possibility shouldexist for flexure and consequent relative displacement of the optical components. This is by no means easy to achieve for an instrument which is to have movement in elevation from horizon to the zenith, and in addition movement in azimuth. Another prerequisite is freedom from vibration in the high winds often experienced Equally any automatic exposure timing mechanisms should function under field conditions. without jarring the instrument. This is of considerable import for exposures lasting more than one night, for clearly the contrast of the images from faint emissions will be enfeebled if the photographic plate is even slightly displaced. In stellar photography it is accepted practice to remove a photographic plate and replace it for the purpose of prolonged exposures making use of a double-slide plateholder. Briefly, this is an arrangement whereby movement of the plate, at the will of the observer, is a means of correcting for shift in the stellar images due to atmospheric turbulence. The relevance to prolonged exposure is that the photographic plate can be re-positioned under control of two micrometer screws working at right angles to one another in the focal plane. So far as the writer is aware no night sky spectrographs have been fitted with a double-slide plateholder. It is certainly worth consideration for a spectrograph involved in both night sky and aurora1 programmes, for the instrument may well be committed to a long duration nightglow exposure and therefore could not be used during the occurrence of interesting aurora1 activity. It would be a relatively simple matter with a double-slide plateholder to remove the nightglow plate, make the aurora1 exposure on a different plate and then replace the nightglow plate and continue the exposure. In 1950 Vegard and Tonsberg w installed a much improved spectrograph at the Aurora1 Observatory at Tromss. The camera operated at f/1*2, with an aperture of 178 mm. The The dispersion collimator and the two-prism dispersing unit had corresponding dimensions. was 41 A/mm at 4000 A, 100 A/ mm at 5000 A and 240 A/mm at 6500 A. The entire optics were mounted inside a well-insulated and thermostatted wooden box, capable of rotation about a vertical axis. An interesting system was used to feed the light into the instrument. A plane mirror was placed in front of the slit at 45” to the collimator axis. The mirror could be turned around a horizontal axis and by means of a condenser lens focused on the slit the is spectrograph could in effect look at any part of the sky. Although such an arrangement undoubtedly a convenient one it is open to criticism on the grounds of the light loss inevitable on reflection at the auxiliary mirror. This loss could amount to 5-10 per cent in practice. Certainly it is desirable to use as few reflecting surfaces, or indeed optical components, as possible in work of this kind, for each one has its attendant reflection or absorption losses. A much faster spectrograph (f/0*65) made by Cojan(i”) was taken to Tromso in 1951. It was intended for studying the more rapid variations of spectral intensity distribution in aurora, as well as for nightglow observations. To give some idea of its capabilities, the nightglow sodium emission was photographed with measurable density in lo-20 min, and during twilight the enhanced D line required 2-4 min exposure. The greater light-gathering power made it possible to follow the twilight sodium emission for longer periods into the night. The same spectrograph also showed the enhancement of H, with increasing altitude in aurora(ll).

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Although the Cojan instrument has nearly 3.5 times the speed photographically of the f/l*2 spectrograph, it had only about half the dispersion. For instance at 4000 A the dispersion was 77 A/mm, 224 ~~rnrn at 5000 A and at 6500 A it was 515 A/mm. The diameter of the f/O.65 objective being 246 mm., the problems concerning the correction of the various aberrations were severe. Clearly there was no need for an achromatic objective, This eased the situation somewhat, as the availability of a variety of glass meant that reasonable correction could be achieved over the spectral range. Generally speaking, curvature of the field is perhaps the most troublesome remaining aberration after all possible care has been taken to eliminate spherical aberration and coma. The curvature of the focal surface limits the spectrum range acceptable by a flat photographic plate. It is unwise to use film, rather than photographic plates, in the focal plane, for there is always the possibility of distortion of the emulsion during the ensuing chemical treatment. The length of the spectrum in the Cojan spectrograph was 12.5 mm over the range 4340-7682 A, the field being flat to within iO.01 mm. It is interesting to note, as Boutry (r2) has pointed out, there is an analogy between a spectrograph camera objective of large aperture and a microscope objective, for after all apart from the colour correction their performance requirements are similar. A single glass prism of the large dimensions involved (240 mm height, and face length 350 mm) would have given far too much absorption in the short wavelength region. Accordingly, the simple expedient was taken of using two 50” prisms of a less dense flint glass thereby gaining greater transparency in the blue. In themselves such large prisms are noteworthy achievements, for it is extremely difficult to obtain glass melts of sufficient quality to provide the necessary blemish free blocks. The collimator lens produced no special difficulties; it was an achromat of 1350 mm focal length working at f/5.6, giving a reduction from slit to photographic plate of nearly 8.5, since the focal length of the camera objective was 160 mm. To study intensity variations of the stronger aurora1 lines it is quite practical to use small single prism spectrographs with a high speed camera lens. For this purpose Harango3) mentions the 50 mm focal length, f/O.95 Tachon objective, made by Astrogesellschaft, Berlin. Exposure times as short as one minute are possible on intense activity. So far as a spectrograph is concerned the night sky and aurora are luminous surfaces of appreciable angular extent. Consequently the light-gathering power of the camera objective determines the sensitivity of the instrument, for the collimator will always be fully illuminated by the incoming radiation. In 1931 a most important step forward in the design of high speed optics was announced by Schmidt (la). Instead of trying to correct the field produced by a parabolic mirror, he considered afresh a spherical mirror. Since spherical aberration affects the optical performance of a spherical mirror, it seemed logical to transfer the parabolizing process, so to speak, to its centre of curvature where every radius can claim consideration as an optical axis. A thin glass plate placed at the centre of curvature, normal to the central axis, and profiled SO as to cancel spherical aberration will correct the images with a high degree of accuracy over a wide field. Fig. 1 shows a classical Schmidt camera. The obvious disadvantage of the arrangement is the curvature of the focal surface; the use of film is imperative unless afield flattening lens is introduced just in front of the focal surface. Meinel(ls) produced an f/O.8 field flattened spectrograph camera in 1948 at Lick Observatory. Acceptable definition covered a plate size of about 17 mm z: 4 mm, the focal length being 100 mm. Unfortunately a field flattening lens placed just in front of the focal plane results in spherical aberration and coma. To overcome these troublesome effects, or at least to reduce them to

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permissible limits, two modifications are necessary. Firstly the field flattening lens should be mounted a small distance from the focal plane. This minimizes the spherical aberration. In the Meinel camera, for instance, the air space between the focal plane and the field-flattener amounted to approximately the thickness of the latter. Moreover, by locating the Schmidt corrector plate not at the centre of curvature, but nearer the focal surface, the coma can be made insignificant. The corrector plate in Meinel’s design was shifted about 20 per cent of the radius of the primary mirror. By curving the plano-convex field flattener lens appropriately, the astigmatic effect was rendered negligible.

Focal

surface

FIG. 1. CLASSICAL SCHMIDTCAMERA.

For airglow and aurora1 studies during the International Geophysical Year of 1957-58 the Cambridge Air Force Research Centre, Massachusetts, U.S.A., arranged for the production of several grating spectrographs fitted with Meinel cameras. The optical arrangement is shown in Fig. 2. The dispersion (2nd order) was 83 &mm. The success of this type of spectrograph, apart from the camera, is a result of the considerable developments in diffracting grating manufacture. The ability to rule an appreciable width is most important; a grating width at least equal to the camera aperture is clearly desirable for minimum light loss in the system. Grating ruling engines of the requisite accuracy have been developed at a Field

Meinel

flattener

came

ractmn ting llimotor

mrror Entrance

Fz.2.

slit

GRATING SPECTROGRAPH.

few Iaboratories, and are capable of producing high quality gratings with 150 mm or more of ruling. Modern diffraction gratings are usually blazed, that is to say the spectral energy is concentrated in a given wavelength region by controlling the contour of the individual rulings or grooves. For the spectrograph under discussion the grating had a ruled area of 127 mm in length and 102 mm width, with 600 grooves per mm. The blaze angle was 19”16’, which corresponded to a wavelength of 5500 A in the second order. One of the difficulties in the use of diffraction gratings is the possible overlap of successive orders. For the case under discussion the second order at 5500 A overlapped the third order in the blue at about 3600 A. As there are a number of strong emissions in the aurora1 tj

680

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spectrum in both these regions it was necessary to block out the blue with a filter. Alternatively a photographic emulsion must be chosen with an appropriately sharp cut-off; for example Kodak 103a-0 plates have little sensitivity above 5000 A and so could be used to record the third order lines without a filter. The mirror collimator increased the potential scope of the instrument in that it could be used in the ultra-violet (the camera corrector plate was of quartz) and the infra-red as well as in the visible region. By virtue of the grating as dispersing element the dispersion was nearly uniform throughout the range. It is perhaps worthwhile to point out that a fairly long focal length collimator, 724 mm in this case, enabled the use of a wider slit than would otherwise have been possible without loss of resolution. Typical slit-widths were 200 ,u for aurora1 activity and 500 ,u for the nightglow. Even though the green aurora1 line from faint (class I) glows was recorded in ten minutes, nightglow spectra containing more detail than just the salient features required exposures of around 8-10 hr. An instrument of this type is quite well suited for investigating enhanced twilight phenomena, exposure times of 10 min producing measurable spectra. It also enables intensity variation studies of the stronger aurora1 emissions, along with wavelength determinations for both the aurora and nightglow. An auxiliary lens placed in front of the entrance slit gives the possibility of investigating the radiations from isolated features of the more intense aurora1 activity. On most occasions, however, the instrument would be pointed directly at the display, without the use of an auxiliary lens, selecting via the collimator slit a particular region. In this manner the radiation from an area of sky 14” in extent is accepted by the collimator. A larger version of this instrument has been in use in Canadau6) for several years. It employs a plane grating of 600 lines/mm and size 8 in. x 8 in. The f/O=8 Schmidt camera is of 9 in. aperture and the linear dispersion 60 A/mm. It is clearly important that all light-losses in the spectrograph should be kept to a minimum, for the overall efficiency, as we shall see later. is not very high. Mirror optics, whereever possible, are a step in the right direction as absorption losses are then at a minimum. The instrument shown in Fig. 2 could be improved in that a more efficient arrangement for the collimator is to mount the slit off-axis and use a long focal length collimating mirror. The corrector plate and field-fattener surfaces alone lead to reflection losses of some 16 per cent of the incident light. By suitable anti-reflection coatings these can be reduced to about 5 per cent. Light is also lost, due to diffraction effects, at the entrance slit. The slit-width is dictated primarily by the grain size and sensitivity of the photographic emulsion. Too wide a slit means that best use is not made of the available resolving power, which is determined by the nature of the dispersing element (prism or diffraction grating). (b) The slitless spectrograph The possible use of a slitless spectrograph at once suggests itself for the brighter aurora1 lines and bands (e.g. 5577,6300,4278 and 3914 A). Under such circumstances Lebedinsky’l’) has obtained sufficiently dense images from bright aurora with exposures of about one minute. Slit type spectrographs of identical focal ratio and emulsion sensitivity require at The arrangement is quite simple least ten times the exposure for the same strong emissions. in that the sky is photographed by a prismatic camera on reflection in a convex mirror. The dispersion of the prism and the distance from the mirror to the camera are chosen so that the monochromatic images of the mirror in the light of the emissions (e.g. 5577,630O A) are

INSTRUMENTAL

OPTICS AND

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just not superimposed. The convex mirror can be masked to leave a very wide slit, parallel to the refracting edge of the prism along its diameter, as shown in Fig. (3a). Such an arrangement would reflect a zone of around 30”-50” in width, intersecting the sky along a complete vertical. Alternatively if the width of the masking slit is reduced until it is several millimeters wide, As the slitless (or strictly speaking, wide-slit) spectrograph becomse a 180” slit spectrograph. before, the exposures are of the order of a minute or so. For convenience, an auxiliary plane Photographic

plate

+ Camera

lens

Prism

(0)

FIG. 3a. SLITLESSSPECTROGRAPH

FIG. 3b. 180" SLITSPECTROGRAPH

mirror located as in Fig. (3b) serves to reduce the overall length of the instrument, which gives a convex mirror to prism distance of 2.5 m otherwise is considerable. Lebedinskyo’) for an f/O.57 camera and two 60” prism spectrographs, and 4 m for a single prism, f/l*25 instrument. Certainly the slitless spectrograph has great potentialities for intensity variation studies of the prominent aurora1 radiations. (c) The eficiency of a spectrograph At this stage it is worthwhile to consider in detail the several factors influencing the efficiency of a spectrograph. Clearly the entrance slit dimensions determine the amount of light passed to the instrument. The wider the slit then the more efficient the arrangement. However, to prevent the wide slit leading to the formation of too broad images in the focal plane a camera objective of short focal length is essential. That is to say its numerical aperture will be large. If dS is the area of the entrance slit, B the brightness per unit area of the entrance slit, then the incident flux on the collimator over a small wavelength range dA is dF = rrB . dl . dS . sin2 or, where sin 8i is the numerical aperture of the collimator. Suppose that for the camera objective we have a numerical aperture of sin t$ and an image surface area of dS’. Then the flux forming the image is given by an amount dF’ = n-B’. dA . dS’ . sin2 8,, where B’ is the brightness

per unit area of the image.

682

A. H. JARRETT

From this we see that the illumination of the image is not affected by the aperture of the collimator lens. For a highly sensitive instrument the relative aperture of the camera objective is all important. Now dF’ is related to dF by dF’=

where T is the where T, is the optics, and To Considering same for each. sion being that

T.dF,

transmission factor for the entire spectrograph. In fact T = T, . To. T,, transmission factor of the entrance slit, To that of the collimator and camera that of the dispersing element. either a prism or a grating spectrograph, then To will be approximately the Fehrenbach(18) investigated diffraction losses at the entrance slit, the concluthe amount of transmitted light increased rapidly with increasing slit width.

Hence in a slit-type spectrograph the nature of the dispersing element is decisive, in that the higher the value of T, the higher the efficiency. For a prism the diffractional resolving power, assuming the entrance slit is infinitely narrow, is given by

s=

i

$i .(tz -

t1)

where 12is the refractive index of the prism material and t, and tz as indicated in Fig. 4. If we take into account the finite width of the slit then the resolving power is found to be

where a is the width of the incident beam (Fig. 4) and F is angle subtended by the slit at the centre of the collimator lens. Evidently for high dispersion the refracting angle of the prism should be as large as possible, consistent with avoiding any light lost by total reflection at the exit face of the prism. In practice the minimum deviation position is generally used as the transmission is then at a maximum. Commenting on the fact that about 60” is the maximum value for the refracting angle of a prism, Boutry(r9) points out that the angle of incidence at the entrance face giving minimum deviation is more than 55” with a reflectivity at each surface of 10 per cent. If the refracting angle is 70”, then for similar conditions, the reflectivity is more than 25 per cent at each face, which is quite intolerable. Although it is advantageous to make the prism from glass with as large a value of dn/dA as possible, a limitation is imposed by the transmission of the glass for the wavelength range in question. This is particularly apparent at the blue-violet end of the spectrum, where the absorption coefficient of the usual varieties of flint glass being fairly high leads to a reduction of the resolving power. Below a wavelength of 4000 A quartz is more suitable in this

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respect, despite its appreciably lower dispersion. Another factor leading to a diminution in resolving power arises from curvature of the image of the slit. The width of the slit image increases with the obliquity of the incident rays, with the result that the resolving power decreases as the image is traversed from the centre to its extremities. By keeping the slit as short as possible in length this effect is minimized. Incidentally, the brightness at the centre of the image depends on the slit width, insofar as it tends to zero when the slit width is very small yet, because of diffraction effects the image subtense tends to the value n/D radians, where D is the diameter of the camera objective. Considering the effect of the grain size (G) of the photographic emulsion we see that to make full use of the resolving power then G 2f. Ay, where Ay is angular extent of the smallest feature resolved in the image andfis the focal length of the camera lens. It can be shown20 that

where A is the angle of refraction of the prism, when used under minimum deviation conditions. Hence from the expression for G we have the focal length and, since the camera lens diameter is determined by the prism dimension, the focal ratio allowing full use of the resolution given by the prism. One outcome of this is that a compromise is required between maximum resolution and a short exposure time corresponding to a small focal ratio. To try for both at the same time imposes impossible conditions. Fehrenbacht2r) has made an instructive comparison of prism and grating spectrographs 6D along the following lines. The angular dispersion of a grating is - = K . m where K is the &? order of the spectrum and m the number of lines per cm of the grating’s ruled surface. cln The angular dispersion of k sixty degrees prisms is g = k .z . tanI.% where n is the II refractive index of the prism glass and 1 the angle of incidence for minimum deviation. Writing g

= k . Y, we can regard a prism as equivalent to a grating of r lines. From his

results of r approximately 1000 lines/cm for FI,, 4000 lines/cm for H;,, dense flint glass for the 60” prism in each case, and 3000 lines/cm at 3500 A for quartz, we see that the dispersive power is appreciably smaller than for a grating of 6000 lines/cm. Bearing in mind the development of blazed gratings by Wood(22) and the Babcocks(23), whereby some 70 per cent of the incident light is diffracted into one order, this being considerabIy more efficient than a prism train of the same angular dispersion, it is clear that a grating spectrograph will give a greater efficiency than one of the prism type. Another point is that the grating keeps its efficiency irrespective of aperture, whereas the efficiency of a prism train decreases with increase in aperture. Turning now to a discussion of the layout of an efficient spectrograph, a variation of a system proposed by EberP4) is shown in Fig. 5a. The system is an off-axis one, but has the advantage that the loss of light arising from the two optical flats of the on-axis system sketched in Fig. 5b is avoided. In Ebert’s original mounting one spherical mirror served both as collimator and camera objective, and as Fastie (25)has shown it is possible to keep the mirror size down to a reasonable limit by locating the grating midway between the mirror and the plane of the entrance and exit slits. The Ebert layout is particularly suitable

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for infra-red studies where the detector is usually a photocell placed behind the exit slit. Under these conditions the off-axis distortions are tolerable. Shemansky and Valiance Jones(26)describe a method of condensing the light from the exit slit of the spectrograph onto the photocell. An ellipsoidal mirror forms a reduced image on the exit slit, and a field mirror images the grating onto the ellipsoidal mirror.

Off-am parabolic mrrore

Diffraction grating

FIG. 5a. EBERT-TYPE SPECTROGRAPH

The light losses incurred by a system as in Fig. 5b are considerable. Excluding the diffraction grating the four reflecting surfaces will lose altogether about 30 per cent of the incident light. Moreover, the 45” optical Aats obscure the mirrors; a conservative estimate would be by as much as IO per cent in each case. So the total loss of light due to reflection and obscuration may amount to around 50 per cent. Entrance

slit Dkffraction

Assuming a blazed grating of 70 per cent efficiency, we see that the overall transmission of the spectrograph is not more than 35 per cent; for an Ebert instrument of the type depicted in Fig. 5a, the corresponding figure is 60 per cent. For a single mirror Ebert spectrograph the figure is further improved to 65 per cent. Clearly for the highest resolution work an onaxis system is desirable, but where efficient use of the incident light is more important than the utmost in resolution (for example in infra-red studies) then the Ebert mounting is superior. II. SYNOPTIC STUDIES (a) Parallactic photography Perhaps the simplest, at any rate with regard to the equipment required, type of work in this respect is that of parallactic photography of amoral features. The Norwegians, under the direction of the late Carl StGrmer 27 have been particularly prominent in this field. Photographs are taken simultaneously from two or more observing stations, 5 cm focal length objectives at not slower than f/2 being quite suitable. Arising from the diffuse nature

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of aurora1 forms, highly corrected lenses are not essential. Absorption of the short wavelength aurora1 emissions is minimized when the objective has as few components as possible. The Astro-RK tine objective of 5 cm focal length (f/1*25) made by Astrogesellschaft(2s) of Berlin has proved very suitable for aurora1 photography. Exposure duration naturally depends on the sensitivity of the photographic emulsion, but exposures of a few seconds on Class 1 activity yield excellent results with modern plates. It is interesting to compare this with the first successful picture of an aurora (a band with ray structure) taken by Brendel(2s) in 1892. The objective was of 210 mm focal length, with a focal ratio of 1: 3.5; the exposure time required was 7 sec. Several important features have emerged from analysis of the thousands of parallactic photographs amassed over the past half century, perhaps the most interesting being that the lower regions of nearly all aurora1 arcs occur at an altitude of approximately 100 km. The International Geophysical Year of 1957/58 attracted much attention to the scheme, organized at Edinburgh by Paton, for visual observations of the aurora. The results of numerous voluntary observers were examined and plotted on specially prepared maps so as to indicate the extent of the different aurora1 forms over the Earth’s surface. (h) The All-Sky

camera

An extension of the above scheme has been made possible by the development of the In 1947 Gartlein(30) successfully used a combination of convex mirror All-Sky camera. and a camera for photographing aurora1 activity. The first reported continuous use, of an The equipment took the form of a cinematograph All-Sky camera is by Lebedinskyc31). camera mounted on a pedestal so as to look down onto a convex mirror. The reflection of the sky in the convex mirror was photographed giving a 180” field of view apart from the One and a half years later Osterbrock and obstruction caused by the camera supports. taken with a wide-angle (i.e. All-Sky) camera designed Sharpless(32) published photographs by Henyey and Greenstein. This had a similar optical arrangement, with a field of view of around 140”, and consisted of a lens which photographed the image of the sky formed by a spherical mirror. The effective focal ratio was 1: 2, the diameter of the field on the original plates being 21.5 mm. For the purposes of the I.G.Y. several All-Sky cameras, either of a type proposed by Meek(33) or Stoffregen’34) were installed at selected sites throughout northern countries and also in the Antarctic. The Meek construction favoured a tine camera facing downwards onto a convex mirror, whereas the Stoffregen design made use of an auxiliary plane mirror to reflect the light from the sky into the tine camera via an opening in the centre of the primary mirror. By placing a number of cameras about 500 km apart it was possible to obtain a simultaneous picture of the extent of the aurora1 activity; the criterion for spacing the cameras was merely that the pictures from adjacent observations overlapped sufficiently to give continuity. The general mode of operation of these cameras was that they photographed the entire sky every minute, with exposure times ofaround 25 set on 16mm tine film, throughout the hours of darkness. Time markings were printed on the film appropriately. In addition the Stoffregen camera possessed a reference grid comprised of a set of small lamps aligned in a N-S and E-W direction over the camera and ranging in elevation from horizon to the zenith. The lamps were recorded on each film frame and consequently the bearing and elevation of any aurora1 feature could be readily determined. The mirrors were front surfaced and protected with a coating of silicon monoxide. A transparent plastic dome(35) fitting over the entire camera was useful for particularly adverse weather conditions.

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The optical design of an All-Sky camera is straightforward in that the amount of the sky observed is fixed by the curvature of the primary mirror, and the separation of the mirror from the camera is determined by the diameter of the mirror and the field of view of the camera objective. One feature of the system particularly worthy of comment is that point sources are not increased in brightness, as they are with a telescope. Generally very few stars are discernible on the films. This, as Evans(36) has pointed out, arises in the following manner. Suppose we consider the light radiated from a source into a small solid angle SO as to just fill the camera Photographic

T

Camera

Interference Collimator

film

lens

fdter

lens

Convex

mirror

FIG. 6. LOCATIONOFANARROWBANDINTERFERENCEFILTERIN

(After

Hunten

AN ALL-SKYCAMERASYSTEM

(40)).

objective, if no mirror is present. The effect of the convex mirror is to increase the solid angle with the result that only a small amount of the light included within the original small solid angle is taken in by the camera. Of course the size of the image in the plane of the film is reduced by the same ratio and so the brightness remains constant. In the case of a star, below the resolution limit of the lens or photographic emulsion, the image is not reduced in size and so the effective response (i.e. the photographic speed) of the system is diminished. For extended sources (aurora1 activity) the photographic speed of the lens plus mirror system is the same as that of the camera objective by itself. Colour photography has interesting possibilities in the province of intensity comparisons of aurora1 radiations. In 1933 Harangc3’) embarked on a series of filter photographs, exposed simultaneously in the green and violet, of various aurora1 forms and obtained isophotes for some arcs and draperies. Developments over recent years in high speed colour film make possible an appreciable extension of this work. Sandford and Heiser(38) have successfully used an all-sky camera with an f/l.4 objective lens in this connection, and withcolour film of 100 A.S.A. nominal speed rating and making use of a forced developing technique could record with a 2 min exposure an aurora just detectable by a visual observer. An interesting

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point here is that aurora of this intensity are below the threshold of human colour vision. Consequently they appear colourless, although they are usually noted as being greenishHowever, the integrating properties of the colour white in appearance by visual observers. film lead to, as Sandford and Heisert3@ point out, the colour latent in these colourless displays being recorded by the film. Feeble activity, appearing white on visual observation, often produces red, purple, blue or white images on the film. The relative spectral intensities derived from spectrograms, taken at the same time as the all-sky pictures, correspond to the colours on the colour film. Monochromatic photographs of aurora are of considerable interest, for example the study of the aurora1 arcs discovered in 6300 8, radiation by Barbier(3g). The narrow band interference filter has made possible investigations of this type, but it does suffer from the drawback of variation of passband with angle of incidence. The location of a narrow band interference filter in an all-sky system is best arranged so that parallel light is incident upon it. In fact, as Huntent40) states, in setting up the optical design one has the choice of imaging the camera objective or the film on the filter. The latter alternative is generally the more convenient. Fig. 6 is a sketch of a suitable arrangement; it has the advantage that the full diameter of the filter, rather than that corresponding to just the aperture of the camera objective is included, so that the angle of the field at the filter is kept to a mimimum. All that is involved is a judicious choice of area of the filter and solid angle in accordance with the Sinet41) condition. Ramsden(42) describes an all-sky camera, suitable for observing monochromatic emissions, similar to the one designed by Hunten (40). For the brighter features a medium bandwidth interference filter (around 50 A halfwidth) is sufficient and the system works well with the filter placed just in front of the camera lens. Another fruitful field for all-sky camera techniques is in the study of inhomogeneities in the nightglow emissions. Patches about 20” in diameter, which last for several hours, have been reported by Manring and Pettit(13) in certain of the nightglow layers. Roach(44) finds that variations in the intensity of the 5577 A line provide additional evidence for the existence of airglow cells some 2500 km in extent. The field of view of an all-sky camera covers an area of about 1000 km in diameter at a height of 100 km. This makes feasible a useful photographic recording of the. intensity distribution of the nightglow. Of particular interest is the possibility of investigating variations in the sodium layer at twilight over the entire sky. A major disadvantage is inherent in the difficulty of differentiating between the image arising from the selected radiation and the one produced from the continuum transmitted by the filter. Interferometer techniques can be used to overcome this possible confusion. Suppose a wedge interferometer is included instead of the filter in the layout of Fig. 6, then a straight line fringe system is photographed by the camera. The important feature here is that the effective pass band may readily be made very small indeed (e.g. of the order of a few Angstrom units wide), since it is a function of the order of interference; this in turn depends on the separation of the interferometer plates. Increase in the reflecting coefficient serves to reduce the pass band. Alternatively, the all-sky camera can be used in association with a Fabry-Perot interferometer, using an arrangement as in Fig. 7. In this case concentric circular fringes are formed in the final focal plane. Whichever system is used gelatine type filters are of course necessary to isolate the desired wavelength region incident on the interferometer. Admittedly, the information from an all-sky interferometer is restricted to the parts of the image crossed by the fringes, but there is a point to point relationship between a particular

A. H. JARRETT

688

Photogrophkc

Camera

ftlm

tens Fabry - Perot Filter

Collimator

interferometer lens

I

c

FIG.~. ALL-SKY FABKY-PEROTINTERFEROMETER

section of a fringe and the corresponding region of the sky. The straight line fringe system makes for easier examination of the film, especially with a microphotometer. The measurement of the rotational temperature of Nzf in bright aurora is mentioned by Ramsden(42). Points to bear in mind are that in order to resolve the components of the band at 3914 A it is necessary to make use of interference fringes of width between 1 A and 2 A. In addition the spectral range of the interferometer must be large enough to encompass the whole band. Using modern high speed photographic emulsions (up to 4000 A.S.A.) it is reasonable to expect worthwhile results from exposures of a few minutes duration with an all-sky interferometer on the enhanced twilight airglow. III. INTERFEROMETERS

FOR AURORAL

AND AIRGLOW

INVESTIGATIONS

(a) The Fabry-Perot interferometer Interferometry has quite a long association with night-sky problems, for as early as 1923 Babcock(J5) used a Fabry-Perot interferometer to determine the wavelength of the green the aurora1 green radiation. night sky line. Vegard and Harang(*6) similarly investigated McLennan(*‘) based his identification of the green line as a forbidden transition in atomic oxygen on the interferometrically measured wavelength. The earlier work lacked the advantage of having the reflecting coatings of the interferometer plates deposited by the evaporation & ~~~~0 process developed by Ritschl(“8) and (3-g), The thermally evaporated films, usually silver or more rarely aluminium, are Strong vastly superior, in terms of homogeneity and minimum absorption of the incident light, to those formed by chemical deposition processes. More recently, evaporated silver films have been superceded by very low absorption dielectric multilayer coatings of the type developed by Banning@O). From the point of view of the efficiency of a Fabry-Perot interferometer the importance of low absorption in the reflecting coatings cannot be overstressed, since the T ‘2 transparency (7) of the interferometer is given by 7 = where A and T are the i W-A i absorption and transmission coefficients respectively of the coatings; clearly it is desirable for A to be as small as possible.

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Rapid progress in multilayer coating techniques since 1949 has led to a renewed interest in the earlier experiments of Babcock. Jarrett t51)has recorded Fabry-Perot fringes from the 5577 radiation with exposures of 20 min from faint aurora1 glows (Class I), using a multilayer coated interferometer and gelatine filter placed in front of an f/2 camera. As mentioned in the previous section, the modern dielectric coated interferometer and narrow band interference filter have interesting possibilities used together with an all-sky camera. The aim of recent interferometric studies (52)of the airglow and aurora has been that of temperature determination from measurements of the fringe profiles. Experimentally the Interferometer

Photomultiplier

problem reduces to a determination of the fringe halfwidth in terms of the order separation of successive fringes. Assuming a Doppler broadening the temperature(53) is given by h = 7.16 x 1O-7 . ~(~~~)~ where T is the temperature in degrees absolute, M the atomic weight of the element, and h is the half-intensity width. The temperature so derived refers to the region of the sky corresponding to the particular section of the image of the interference fringe. Photographic recording of the fringe system has the obvious advantages of simplicity and that the plate can be microphotometered at leisure. However, the considerable improvement in the sensitivity of photomultiplier tubes during the early 1950”s made it possible to scan Fabry-Perot interference fringes photoelectrically. Armstrong(“2&) arranged a photomultiplier in conjunction with a Fabry-Perot interferometer in a system as shown in Fig. 8. His equipment was based on a method suggested by Jacquinot and Dufour(54), whereby the fringes are imaged on a mask having a hole in its centre and a series of concentric circular slits arranged to coincide with successive orders. The light transmitted by the mask falls on the photomultiplier tube. If the order of interference is varied, by altering the density of the air between the interferometer plates, then the circular fringes expand or contract radially. That is to say the fringes progress across the slits. A convenient way of effecting the necessary change in air density is to raise the pressure in the cylinder containing the interferometerand then plot the photomultiplier current against pressure as the air is allowed to leak away. A later version made use of ratemeter techniques, that is the photomultiplier was used with a pulse counting circuit. Karandikar(~~b) used a similar arrangement and arrived at the same conclusion as to the success of the method, namely that for temperature measurements over stationary regions of aurora1 activity or for surveys across a quiet homogeneous arc it is quite promising. Comparing the work of Wark and Stone(55), who used photographic recording, with the photoelectric methods of Armstrong and. also Karandikar, we see that whereas exposures of several hours on the nightglow were necessary with photographic plates only 20-30 min

690

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H. JARRETT

sufficed for the latter. The main limitation to speed of response with photoelectric recording is that the rate of change of pressure in the interferometer cylinder must be small to avoid adiabatic effects in the air gap which would affect the optical path difference. One way around this difficulty is to use a gas of high refractive index, since (,u - 1) will then be greater than that for air. Although pressure scanning has been adopted by most workers in this field other schemes are possible. For instance Dupeyrat W) has shown that a piezo electric spacer can be used in the interferometer and a voltage applied to give a small displacement. Slow rotation of the interferometer about the optical axis of the system has been suggested by Jaffec5$). Tolansky and Bradley@) have constructed a photoelectric spectrometer in which the spectral profile is scanned by mechanical movement of one of the plates of a Fabry-Perot interferometer. The movement is actuated by a moving coil vibrator at frequencies of up to 1000 c/s; this makes the instrument very suitable for rapid measurements of spectral intensity. Recent achievements in artificial satellite development have provided added interest to pulse-counting possibilities, for clearly we have here a way of obtaining information that can be readily handled by a telemetry system; all that is required is a digital recording of fringe intensity against order of interference. It is worth bearing in mind, as Karandikar’““b) points out, that in the photoelectric scanning of a fringe system we replace the traditional use of the narrow selecting slit of a microphotometer, to obtain a density trace along a diameter of the photographed fringe system, by the direct recording of the light passing through the circular slits in the photomultiplier mask. Considering the factors governing the sensitivity of a scanning interferometer we can take as a commencing point that the current at the photocathode i is given by i=-&.A.Q.T.E.F,

where I is the photon intensity cm-* se& sterad-l, A is the area of the effective aperture of the instrument in cm2, L2 is the solid angle subtended by the field of view of the instrument, T is the transmission factor of the entire optical system, E is the photoelectric efficiency of the photocathode for the spectral band under consideration, and F is the conversion factor we may rewrite the above expression as i = q~ x for electrons/set to amperes. Alternatively net gain factor, where p; is the flux incident on the interferometer photocathode. It is important to realise that the field of the instrument is very much smaller than that generally used in normal airglow photometers: the field is of course defined by the central diaphragm. The simplest possible mask is one with only a central opening. If the radius of this is made equal to the width at half intensity of the central fringe then it can be shown(61) that the system has a resolving power 0.7 that of the interferometer considered separately. Moreover, this arrangement makes for a reasonable compromise between the effective resolving power and the light flux transmitted. Now Jacquinot(5g) has shown that the flux p transmitted through a mask aperture of diameter equal to the central fringe halfwidth is given by FE&L.“, 47r where T is the transparency

of the interferometer,

R, and R, is the resolving

power

of the

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interferometer. Taking realistic values, such as T = 0.9 (for multilayer coatings of O-80 reelecting coefficient and 0.01 absorption coefficient), R0 = 400,000 (for an order of interference of 30,~0), A = 78.5 cm2 (for interferometer plates 10 cm clear aperture), and I = 2.5 x 108 photons cme2 see-l sterad-l for the average 5577 A airglow@), then we have q~ = 17,300 photons/s. We can arrive at a reasonable value for the net gain factor as follows : supposing about 10 circular rings (it would be very difficult to use any more than this number in practice) are used in the photomultiplier mask in addition to a central hoie, then a gain in light fluxof say3 by a single central hole. There(see Karalldikar(5zb)) may b e achieved over that transmitted fore we could expect about 52,000 photons/s from a mask as just described. We must now take into account losses arising from reflection at the air-glass surfaces in the complete optical system (see Fig. 8). Twelve are involved, so with all surfaces anti-reflection coated a transmission of (O-98 )I2 equivalent to 78 per cent of the incident radiation is the maximum possible. Assuming that the interference filter is of the narrow band multilayer type then its transmission is unlikely to be more than 60 per cent for a bandwidth of 20 A. For a typical highly sensitive tri-alkali photomultiplier the quantum efficiency is perhaps 8 per cent. Consequently the flux at the photocathode is equivalent very nearly to 3.12 x lo-l6 amp, since F = 1.6 x lo-is. From information concerning the exponential extinction coefficient due to the atmosphere it is reasonable to assume, as Bates and Dalgarno f621have pointed out, that the intensity at a zenith distance of 75” is twice that at the zenith, This would give a photocathode current of about 6 x lo-l6 amp for observations made at such a zenith distance. Certainly this is a very small current, but with a photomultiplier gain of 5 x lo6 as typical, the anode current is 30 x lo-lo amp. This could readily be increased by a factor of lo5 in the amplifier, giving an output of 300 p amp for recording purposes, which is more than adequate. A troublesome aspect of photomultiplier tubes is the thermionic emission or dark current from the photo-sensitive cathode. It may easily amount to the order of lo-l6 amp at room temperature and generally it is essential to arrange for cooling to solid carbon dioxide or liquid air temperatures. This can be inconvenient, especially under field conditions.’ It is advantageous to include a discriminator in the output circuit of the photomultiplier to reject the low amplitude dark current pulses which constitute background noise. Since one tenth of the output calculated above for the 5577 A nightglow radiation is sufficient for recording purposes we could expect to obtain a measurable output from radiations of intensity not less than 25 x lo6 photons cm-2 se& sterad-l (i.e. 25 R). According to emission rates quoted by Chamberlain (60) this would include the [OIg 6300 A and Na 5893 A niglltglow radiations. II, is barely possible as its emission rate is at most only 20 R in the direction of the zenith; at the minimum emission rate of 5 R it certainiy is not detectable. Twilight studies of the 6300 and 5893 radiations are possible, along with the lithium 6708 and calcium 3933 emissions. Studies of these other wavelengths presuppose photosensitive cathodes of comparable efficiency to that of 8 per cent taken for the 5577 line. A fascinating possibility of the techniques of photoelectrically detecting Fabry-Perot fringes lies in producing the line profile directly on a cathode ray oscilloscope. If the radiation is sufficiently intense (from strong aurora1 activity) for the variation in photocurrent output to be recorded in a short time interval, then rapid scanning of the fringe system would allow the multiplier tube output, with suitable amplification, to be displayed on an oscilloscope screen. The scanning rate would have to be such as to encompass several fringes per

692

A. H. JARRETT

second; one suggestion for this is the action of a cam-operated pistontSzb) arranged to produce time-linear pressure changes in the Ctalon gap. Another line of attack on the temperature problem hinges on the visibility of the fringes. The visibility Vis defined as equal to (I,,, - ZtniJ/(Zmas+ Zmin)where ZmaX is the intensity at the peak of a fringe and _Imin the intensity at the minimum between two fringes. Assuming the fringes are Doppler broadened, SahatG) developed the expression

wherep is the order of interference, c the velocity of light, k is Bol~mann’s constant, T is the absolute temperature, M the atomic weight of the emitting atom and MB the mass of the hydrogen atom, and R the reflection coefficient of the Btalon plates. For determinations of this kind a much lower resolving power willSsuffice than is essential for accurate profile plotting. The consequent reduction in reflection coefficient’for the &talon plates leads to an enhancement of the transmitted flux, which might well prove decisive for a feeble nightglow emission. Armstrong(@) has shown that if the intensity of the background sky light is neglected in this type of investigation then spuriously high temperatures are given. Roach and Pettit(65) have obtained a representative value for the total background light of the night sky and found it equivalent to 500 tenth magnitude stars per square degree in regions remote from concentrationsof stars. Photographically the result is that the fringes are superimposed upona general fogging of the plate. The effect of the fog must be determined and then avalue for the temperature can be deduced from the corrected contrast. The stronger the intensity of the emission (in an aurora for example) then the smaller are the errors arising from the background light. Another issue ensuing from the use of low-reflecting coefficient &talon plates is that the instrumental width becomes an appreciable fraction of the observed fringe width, with the result that the Doppler broadening is subdued. In experimental work of this kind it is most important to ensure that the linear diameters of the fringe images are sufficiently large to avoid any troublesome effects arising from the grain size of the photographic emulsion during profile or intensity measurements. (b) The Michelson ~~~erfero~e~er The interferometer conceived by Mi~helsont66) in 1881 has been neglected over the years, so far as studies of upper atmosphere phenomena were concerned, until it was pointed out by Fellgett(67) in 1958 that any two-beam interferometer together with Fourier transformation has the great advantage of simultaneous observations. MichelsonV*) pioneering research on the structure of atomic lines rested on visibility measurements of the interference fringes. Alone the interferogram will not yield explicit information relevant to the spectrum from which it is formed, but since the fringes are physically determined when the spectral energy distribution and the characteristics of the interferometer are decided, then conversely information about the spectral features can be derived from knowledge of the fringes and interferometer. Information concerning the spectrum is inherent in the interferogram as a Fourier transform. Figure 9 shows the optical arrangement of a Michelson interferometer suitable for night sky work. Perhaps the most difficult technical feature is to ensure that the moving reflecting surface always remains strictly parallel to its initial position. Mounting the mirror carriage on optically polished sliding ways is one way of accomplishing this.

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The resolving power is theoretically unlimited, being given by R = 2~ . D where v is the maximunl frequency observed and D is the distance traversed by the mirror; in practice a resolving power of 5000 or even 20,000 can be achieved. The light-gathering power is determined by the clear aperture of the interferometer mirrors, analogous to the aperture of the optical flats in a Fabry-Perot Ctalon. The optical components must be of very high quality, their faces worked to one fortieth of a wavelength and in addition the beam splitter and compensating plates homogeneous throughout; these stringent conditions limit the aperture in practice to around 10 cm. The Michelson interferometer together with the Fourier transformation technique is admirably suited for studies of the hydroxyl bands around the I.5 ,u region, and provides an

Off - axis

parobok

Beam

splitting Ond CO~Oensofing plates

FIG. 9.

MICHELSON

INTERFEROMETER

FOR NlGHT

SKY

INVESTIGATIONS

accurate means of deriving temperatures from the rotational-vibrational theory of the bands. Temperature and intensity variations are well within its scope, for the scanning time is around 15 min. This makes it well suited for diurnal investigations. The conversion from the interferogram to a spectrum involves a digital computer, the final resolution being a function of the number of points to which the Fourier transform is applied. However, the really important point is that when used in the infra-red with photoconductive (e.g. lead sulphide) detectors the signal to noise ratios of the spectra produced are appreciably superior to spectrographic equipment of comparable resolution. This gain stems from the fact that all the components in the spectrum arrive at the detector for the entire duration of the scan; in a conventional scanning spectrograph a spectral feature is recorded for only a fraction of the time taken to expose a particular region of the spectrum. (c) Merits of spectrographic and interferometric methods For similar resolving power, a spectrometer with a Fabry-Perot interferometer as its dispersing element is capable of collecting more of the incident radiant energy than a slittype spectrograph,since the solid angle of the light beam accepted by the etalon is appreciably larger. This of course means that there is more light flux at the detector of a spectrometer incorporating a Fabry-Perot &talon. Apparently then for low intensity sources, such as are met with in night sky emissions the interferometer is superior. However, it must be borne in mind that the Fabry-Perot interferometer is not suitable for investigating extended wavelength regions, for the spectral range of the interferometer {i.e. the waveIength region over which it can be used without overlapping of successive orders of interference) is given by the well-known expression AA= -!! where t is the separation of plates. In fact, for the problems 2W under discussion this amounts to an investigation of one wavelength at a time. In any event,

694

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the interferometer must be preceded by a narrow band filter (effectively acting as a monochromator) so as to exclude all spectral elements outside the range An. Against this we have the simultaneous information over a considerable spectral region provided by a grating spectrograph. Clearly it is essential to study the special requirements posed by an investigation so that a correct assessmentcan be made of the suitability of a particular spectrometer. For instance if the problem is to investigate a given spectral region with a certain resolving power, then the merit of the instrument can only be decided by a comparison of total scanning time with the interferometer and photographic exposure time with the spectrograph. A major criticism concerning interferometric methods lies in the manner in which the information is presented; it does not possess the line image simplicity of a slit-type spectrograph. The reduction of the data provided by a Fabry-Perot fringe system is complex (c.f. Tolansky(6g)), but in its turn is very much simpler than the Michelson interferometer plus Fourier transform technique outlined above. Nevertheless the Michelson interferometer possesses the great advantage that all the spectral elements are investigated simultaneously instead of successively as with the Fabry-Perot instrument. This means that there is a gain in the duration of the measurement amounting to the number of spectral elements investigated. Technically the Fabry-Perot is easier to construct and operate than a Michelson interferometer. Its compact dimensions make for ease in temperature control and moreover there are no worries over the most difficult aspect in operating a Michelson interferometer, namely the accurate movement of one of the mirrors so as to always remain parallel to its initial setting. In many cases photoelectric detection is preferable to the photographic plate; for instance where instantaneous response is desirable. Jacquinot (‘O)has remarked that whilst the Fourier method has possibilities similar to that of the photographic plate it is extremely useful in the region for which no sensitive emulsions exist (e.g. the 1.5 ,u-2.2 ,LL)region of interest to aeronomists. Certainly for the near infra-red, say over a spectral range l-2 p, the Michelson interferometer is a powerful instrument. Even with modest resolving power of 5000, a resolution of 3 A is possible for the l-5 ,LLregion. The significance of this is perhaps best illustrated by comparison with the resolution of an Ebert spectrograph specially constructed for nightthe limit is a spectral slit width of 25-50 8, at 1.5 ,u. The glow infra-red investigations; spectrographic procedure for temperature determinations is to fit a synthetic spectrum derived from rotational theory of the OH bands, with an experimental profile obtained from averthe errors involved are not smaller than f20”K (for a aging several observed spectra; temperature of 270°K). Sources of error are several, noise in the spectra and deviations from the resolution function of the spectrograph being amongst the most prominent. Using a Michelson interferometer it is reasonable to expect an error of only f5”K in temperature determinations from the hydroxyl emission around 1.5 p. Spectrographic studies of the near infra-red out to a wavelength of 1.2 ,u have been made This device originated in a desire to explore the possible with the aid of image convertors. use of electronic techniques as a supplement to photographic methods for increasing the wavelength sensitivity of optical instruments, in particular telescopesc71). Modern photographic emulsions, hypersensitized prior to use, are responsive in the near infra-red as far out as 1.0 p. The essence of the image converter is that the image formed optically on a photosensitive cathode is reproduced on a photographic emulsion, making use of the sensitivity cathode of the emulsion to high speed electrons. Electrons leaving the partially transparent

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are accelerated and focused via an electron lens (see Fig. 10) so as to converge on the photographic plate. It is perhaps appropriate to remind the reader of the action of high speed The exposure is determined by the electronic charge electrons on photographic emulsions. (i.e. the product of the number of electrons per second and the time) reaching the emulsion. There is no reciprocity failure, in that a sufficiently long exposure will lead to a blackening however feeble the stream of electrons. The gain in efficiency over direct photography is considerable. As regards the sensitivity limit the dark current from the photocathode ultimately means that the photographic plate is blackened all over for very long exposures; the standard technique of using a cooled cathode is of help in this respect. Using a caesium-antimony cathode it is possible to make image converters sensitive up to about 1.2 ,u. Krassovskii(72) has carried out work on the infra-red nightglow emission by Electron

lens

Pnotosensltlve /

Electron

Photographic plate

image

FIG. 10.

photographing the spectrum with an image converter, his studies of the hydroxyl emission extending to 1.1 ,u. Bagariazkyand Fedorova (73)have obtained aurora1 spectra up to 11,800 8, with the aid of image converters. Their earlier spectra were of low dispersion, 1200-3000 A/ mm, using a prism spectrograph. Later they improved matters with a f/l*0 grating spectroA/ mm dispersion, 10 A resolving power, and with a range of nearly 1500 A graph of 230-250 in extent capable of being photographed at one time. The presence of helium emission at 10830 A in sunlit aurora1 spectra has been detected by Fedorova and Shefov(74). Vegard(75) has investigated aurora1 temperatures by means of the intensity distribution of the rotational lines in the R branch of the 3914 A and 4278 A N,+ bands. He arrived at the conclusion that the temperature is independent of the altitude of the activity, a result at variance with other evidence which indicates an increase of temperature with height above the 100 km level. Other workers, for example Shepherd and Hunten(76), Valiance Jones and Harrison(77), and Montalbetti (78),all using high dispersion spectrographic equipment found their results in accord with a temperature gradient of about 6” per km. This in turn agrees with the temperature profile as obtained by rocket-borne apparatus. Hunten(7g) suggests that the constant temperatures found by Vegard most likely arise from contamination of the spectra of high altitude aurora with light from the usually brighter lower altitude activity, the contamination being brought about by scattering of the light. The net outcome of this is that the temperature always corresponds to that of the low altitude display. A recent development devoted to the investigation of very faint sources with a spectrometer evolves from the idea that the averaging of a considerable number of short duration spectral scans is, for all practical purposes, equivalent to the information yielded over a long exposure time (i.e. scanning time). For instance low leakage condensers can be used to store the signal from a photodetector, and if the very weak signal is repeated many times the signal to noise ratio builds up in proportion to the square root of the number of scans. Admittedly, the final signal to noise ratio is only slightly better than that obtained from a

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single long exposure time. Such a scheme helps to overcome the limitations imposed by variations in the intensity of the airglow during a long exposure. Hunten@O)describes a condenser memory unit capable of scanning 32 channels in 10 sec. This number of channels is sufficient for looking at one or two isolated lines in a spectrum but is of little use for comprehensive studies of airglow and aurora1 features. Using more sophisticated electronics Broadfoot and Hunten have constructed a 600 channel memory unit. This enables about 1000 A of spectrum to be scanned with a resolution of 5 A, allowing three channels for averaging per spectral feature. The thermal emission from water vapour and carbon dioxide in the lower atmosphere becomes sufficiently strong beyond 2.5 p to overshadow the airglow emissions from higher altitudes. Consequently for spectral studies in the infra-red beyond 2.5 ,u it is necessary to carry out the observations from balloons or high flying aircraft. The problems presented by this thermal radiation and strong thermal emission from parts of the spectrographic equipment itself have been discussed by Valiance Jones@2). IV. PHOTOMETRY

OF THE AIRGLOW AND AURORA

Attempts at visual photometry of the aurora have been very sparse and of course quite impossible for the airglow. The only visual photometer for aurora1 colour measurements known to the author is the one Gadsden(E3) used for colour investigations of the zenith twilight sky. Photometers making use of photoelectric detectors have been used quite extensively, the usual arrangement being an objective lens, preceded by a filter, with the photodetector located in its focal plane. The usefulness of this kind of photometer has been vastly increased by the successful development of interference filters, which are considerably more efficient than their wideband dye type predecessors. To mention some of the problems attacked in this manner, we have the investigations of Roach and Meinel@*) on the height of the 5577, 6300 and 5893 A nightglow emissions by the Van Rhijn method. Similar equipment was used by Huruhata, Tanabe and Nakamuracs5) in their studies of the height of the 5577 A emission. For convenience in this work the photometer is generally set on an altazimuth mounting, and the driving motor so programmed that after finishing a sweep in the vertical at one azimuth the photometer is reset to the initial altitudepositionandthe~imu~changed. In this way automatic coverage of the sky is ensured. The intensity distribution of the airglow emissions has been investigated with similar photometric equipment, for example the work of Manring and Pettit(@) concerning 5577 A. An extensive photometric investigation of the light of the night sky was performed by Barbierta7); his measurements also included the OH bands near 6300 A. Diurnal variation studies of OH emission in the nightglow have and also byRoach(91)and his associates. been made by Elvey cg8),Huruhata(ss), Armstrong( Roach, Pettit and Williams(91)interpreted their photoelectric photometer observations of the night sky over a wavelength range of 0.64 ,u to 1.1 ,u in terms of the height of the OH emitting layer. They give this as being at an altitude of between 50 and 90 km. Armstrong’s@o) photometer, used a Wratten infra-red filter and a caesium-oxide-silver cathode gas filled photoelectric cell, the wavelength range accepted being from 7400 A to 11,300 A. The field covered by the instrument was approximately 6” x 12”. The observations of Huruhata(ss) are most interesting in that he observed a moving pattern in the radiation. Earlier than this Elvey, Swings and Linketg2) reported that they had found evidence of patchiness in the emissions from the night sky, but it was not until the development of the automatic photometers mentioned above that isophotal maps became

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an

accomplished fact. In this connection should be mentioned the work of Roach and Pettit on the NaD lines(s3) and 5577 A(s~~, Photoelectric photometers have also been extensively used in investigations of the variations of the NaD emission in twilight. A comprehensive comparison of various methods has been made by Hunten(s5). The gegenschein and zodiacal light were studied by Roach and Rees(@Qusing a photoelectric photometer of the type used in airglow observations. 3irefringent filters(s7), capable of transmitting very sharp wavebands (of the order of a fraction of an angstrom) have been incorporated into airglow photometers; Koomen, Packer and Tousey(gs) investigated the intensity of the 5577 line in the night and twilight airglow. St. Amand describes a photometer whose response is proportional to the logarithm of the incident intensity and possessing long-period stabihty of calibration by virtue of comparison with a self-luminescent phosphor excited by a carbon 14 source. The measurement of the absolute brightness of the airglow emission is fundamentally a question of photometry at low intensities. The usual procedure is to make use of a diffusing screen of magnesium oxide illuminated by a standard lamp. Alternatively stars can be used as standard point sources. There is, however a likelihood of error from the latter, when employed for surface photometry calibration, arising from the possible lack of uniformity of response of the photosensitive cathode. At the time of writing the only dayglow emission, as distinct from those of the nightglow, detected by sea level observations is that of sodium. Blamont and DonahueuoO) found it to have an unexpectedly strong intensity, some 30 kilorayleighs towards the zenith. Their photometer was a magnetic scanning one and used a sodium vapour absorption cell. Advantage was taken of the fact that the NaD lines are resonance lines and the cell served to scatter the light from the sky into a p~otomultiplier tube. Discrimination between the portion of the output signal having its origin in the dayglow sodium and that resulting from Rayleigh scattering was achieved by periodic application of a magnetic field to the cell, in such a manner as to displace the Zeeman component of the lines in the cell from the atmospheric sodium lines. 1. D. R. BATES,Physics of the Upper Atmosphere, (ed. J. A. RatclXe) Academic Press, New York &London, p. 219,196O. 2. F. E. ROACHand P. M. JAMNICK,Sky Telex., 17, 8, (1958). Physics of the Aurora and Airglow, Academic Press, New York & London, p. 571, 3. J. W. CHAMBERLAIN, 1961. 4. V. M. SUPHER,Astrophys J. 49,266, (1919). 5. LORD RAYLEIGH,Pro&. Roy. Sm. AlOO, 367, (1922). 6. A. J. ANGSTROM,Ann. Phys., Lpz. 137,161, 1869. 7. K. KAYSER,Handbuch der Spektroskopie, 5, 47, Hirzel, Leipzig, (1910). 8. D. M. HUNTEN,J. Atmos. Terr. Phys., 7, 141, (1955). 9. L. VEGAR~Iand E. TCINSBEKG, Geofys. Publ., 18, 3, (1952). 10. J. COJAN,Ann. astrophys., 10, 1947. 11. L. VEGARD,Natlcre, Land., 170, 536, 1952. 12. G. A. BOUTRY, ~nstr~rnent~l optics, (trans. R. Auerbach), Hilger and Watts, London, p. 530, 1961. 13. L. HARANG, The Aurorae, Chapman & Hall, London, p. 64, 1951. 14. B. SCHMIDT,Mitt. Hamburg. Sternw. 7, 15, (1932). 15. A. B. MEINEL,Mitt. Hamburg. Sternw. 23, 15, (1955). 16. D. H. MCPHERSON & A. VALLANCEJONES,J. Atmos. Terr. Phys., 17, 302, (1960). 17. A. I. LEBEDINSKY, The Airglow and the Aurorae, (ed. E. B. Armstrong & A. Dalgarno), Pergamon Press, London & New York, p. 388, (1956).

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