A SURVEY OF RADIAL VELOCITIESin theZODIACAL DUST CLOUD

Brian Harold May
Astrophysics Group
Department of Physics
Imperial College
LondonThesis submitted for the Degree of Doctor of Philosophy toImperial College of Science, Technology and Medicine
London· 2007 ·2

Abstract
This thesis documents the building of a pressure-scanned Fabry-Perot Spectrometer,
equipped with a photomultiplier and pulse-counting electronics, and its deployment at
the Observatorio del Teide at Izaña in Tenerife, at an altitude of 7,700 feet (2567 m), forthe purpose of recording high-resolution spectra of the Zodiacal Light. The aim was toachieve the first systematic mapping of the MgI absorption line in the Night Sky, as afunction of position in heliocentric coordinates, covering especially the plane of theecliptic, for a wide variety of elongations from the Sun. More than 250 scans of bothmorning and evening Zodiacal Light were obtained, in two observing periods –September-October 1971, and April 1972. The scans, as expected, showed profiles
modified by components variously Doppler-shifted with respect to the unshifted shapeseen in daylight. Unexpectedly, MgI emission was also discovered. These observations
covered for the first time a span of elongations from 25º East, through 180º (theGegenschein), to 27º West, and recorded average shifts of up to six tenths of an
angstrom, corresponding to a maximum radial velocity relative to the Earth of about 40
km/s. The set of spectra obtained is in this thesis compared with predictions made froma number of different models of a dust cloud, assuming various distributions of dust
density as a function of position and particle size, and differing assumptions about their
speed and direction. The observations fit predominantly prograde models fairly well,
but show a morning-evening asymmetry, different in the two observing periods.
Models are investigated containing various components, including prograde andretrograde orbiting dust around the Sun, a drift of interstellar material though the Solar
System, and light from distant emitting matter.

The implications for possible

asymmetries of the Zodiacal Cloud are discussed. Other researches on the Zodiacal
Dust Cloud, before, during, and after my observations, are reviewed, including recent
insights into its structure, orientation, and evolution, up to the present day, and myobservations are evaluated in this context. Period of study, 1970-2007.

3

Contents
Abstract2

Declaration and Copyright7

Dedication

8

Preface and Acknowledgement

9

List of Figures

14

Abbreviations and Acronyms

17

1 Introduction

18

1.1 The Zodiacal Light – the phenomenon . . . . . . . . . . . . . . . . . . . . . . . . . . . .

18

1.2 Review of observations of the ZL prior to 1970....................

19

1.2.1 Earliest views . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

19

1.2.2 Evidence as to the scattering medium

26.......................

1.2.3 Measurable quantities in ZL observations...................

27

1.2.4 Physical parameters and models of the ZL..................

29

1.2.5 Doppler investigations of the ZL prior to this study...........

1.3 Review of the period of my observations of the ZL, 1971-197433.......

41...........

42

1.4.1 Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

42

1.4.2 Infrared surveys . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

46

1.4.2.1 Modelling based on IRAS data . . . . . . . . . . . . . . . . . . . . .

46

1.4.2.2 ISO . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

54

1.4.2.3 COBE . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

57

1.4.3 Collected samples and in situ measurements . . . . . . . . . . . . . . . . . .

58

1.4.4 Interstellar dust . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

60

1.4.5 Doppler measurements and models since 1974 . . . . . . . . . . . . . . . .

61

1.4 Review of recent developments in ZL research, 1974-2007

2 Preparations and experimental details 1971-1974

64

2.1 Choice of wavelength for the observations . . . . . . . . . . . . . . . . . . . . . . . . .

64

2.2 Choice of resolving power . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

66

2.3 The site . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 66
2.4 The ZL Building . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

68

2.5 Instrumental details . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

69

4

2.5.1 The coelostat and sky area accessible . . . . . . . . . . . . . . . . . . . . . . . .

69

2.5.2 The Fabry-Perot interferometer (F-P) . . . . . . . . . . . . . . . . . . . . . . . .

72

2.6 New features of the instrument . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

76

2.6.1

Fore-optics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

76

2.6.2

Field lens . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

76

2.6.3

Beam splitter for compensation channel . . . . . . . . . . . . . . . . . . . . . . 77

2.6.4

Construction of housing . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

2.6.5

Fabry-Perot theory, and the finesse of a practical interferometer . . . 79

78

2.6.5.1 Reflection finesse . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 80
2.6.5.2 Aperture finesse . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83
2.6.5.3 Defect finesse . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 85
2.6.5.4 Coating non-uniformities . . . . . . . . . . . . . . . . . . . . . . . . . .

87

2.6.5.5 The choice of plates . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

89

2.6.6 Intensity calibration box . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

92

2.6.7 Wavelength calibration assembly . . . . . . . . . . . . . . . . . . . . . . . . . . . .

93

2.6.8 Control electronics......................................

94.........................................

99

2.6.9 Output devices

2.7 Optimisation of the system
2.7.1 Optics.....................................

99................................................

99

2.7.2 Counting electronics. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 100

2.7.3 The system as a whole. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 100

2.8 How the observations were made. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 101

3 Reduction of the Data105

3.1 Intensity calibration. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 105

3.1.1 Method and calculations. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 105

3.1.2 Accuracy of intensity calibration. . . . . . . . . . . . . . . . . . . . . . . . . . 109

3.2 Wavelength calibration . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110
3.2.1 Determination of the datum line position
3.2.2 Establishment of wavelength scale
3.3 Treatment of a single scan

. . . . . . . . . . . . . . . . . . . . 110. . . . . . . . . . . . . . . . . . . . . . . . . 111. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 114

3.4 Daylight correction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 117
3.5 MgI emission. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 119

5

3.5.1 Discovery . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 119
3.5.2

Wavelength of emission . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 121

3.5.3

Intensity of the emission . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

3.5.4

Variation of daily average emission . . . . . . . . . . . . . . . . . . . . . . . . 123

3.5.5

Variation of emission throughout the night . . . . . . . . . . . . . . . . . . . 124

3.5.6

Final comments on MgI emission . . . . . . . . . . . . . . . . . . . . . . . . . . 136

121

3.6 Averaging . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 127
3.7 Polynomial fits . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

128

3.7.1

General . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

128

3.7.2

Computation of the minimum point . . . . . . . . . . . . . . . . . . . . . . . . 131

3.7.3

Computation of shift by the ‘inflexion points’ method . . . . . . . . . .

133

3.7.4

Assessment of these methods . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

135

3.8 Gaussian Curve Fitting . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

136

3.8.1

Computation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

136

3.8.2

Experimental case for the use of Gaussian fits . . . . . . . . . . . . . . . .

136

3.8.3

Theoretical justification for the use of Gaussian fits . . . . . . . . . . . .

137

4 Interpretation of results in terms of physical models

140

4.1 Gallery of fully reduced scans . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

140

4.2 General comments on the data.................................

151

4.3 Information obtainable from the Gaussian fits . . . . . . . . . . . . . . . . . . . . .

152

4.4 Accuracy of the fits . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

156

4.5 Evidence from line-widths and depths . . . . . . . . . . . . . . . . . . . . . . . . . . . .

157

4.6 Analysis of the wavelength versus elongation data . . . . . . . . . . . . . . . . . .

159

4.7 Forces experienced by particles in the Zodiacal Cloud . . . . . . . . . . . . . . .

159

4.8 Rotating dust cloud models . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

160

4.9 Other considerations in rotating models

163

4.10 Geocentric Dust Cloud (GDC) model
4.11 A uniform dust flow model..........................

. . . . . . . . . . . . . . . . . . . . . . . . . . . . 164. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 166

4.11.1 Uniform flow theory . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

167

4.11.2 Comparison of continuous flow model with experiment . . . . . . .

175

4.12 Estimation of the effect of a contribution to the Zodiacal Light fromsources distant from the Solar System . . . . . . . . . . . . . . . . . . . . . . . . . . . .

176

4.12.1 Computation

176.........................................

6

4.12.2 Comparison of distant emitting model with the data . . . . . . . . . . .

180

4.13 Further analysis of the data by comparison with new radial velocity data
and model theory available in 2007 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

181

4.13.1 Observations of Fried 1977 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 181
4.13.2 Observations of East and Reay 1984 . . . . . . . . . . . . . . . . . . . . . . . 182
4.13.3 The WHAM observations, and theoretical model predictions . . . . 183
4.14 Binning of the HMR data . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 186
4.15 Variations with time

. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 188

4.16 Summary of implications of the 1970-4 work, and pointers for the future190

5 Current developments and future plans

192

5.1 Current Zodiacal Light research192................................

5.2 Objectives of our new work in Zodiacal Light.....................

195

5.2.1

Doppler measurements on scattered light at optical wavelengths . .

195

5.2.2

Doppler shifts of emission lines in the infrared . . . . . . . . . . . . . . .

197

5.3 Concluding remarks . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

198

5.3.1

Summary............................................

198

5.3.2

Visions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

199

Appendix 1 Program ZLPLT207

Appendix 2 Programs LSQFIT, GAUSSN

210

Appendix 3 Programs POLYPT, FLEXPT, MINIPT, MINIPS, GAUSPT

220

Appendix 4 Programs STREAM, CONVL

224

Appendix 5 Hicks, May and Reay NATURE paper, 1972

228

Appendix 6 Hicks, May and Reay MNRAS paper, 1974

233

Appendix 7 The Airy Summation

244

References

247

7

Declaration and CopyrightThis thesis is my own work, except where explicitly indicated in the text, and copyright© Brian May 2007.
With reference to my 1970s publications on this work, in Nature and MNRAS (seeAppendices 5 and 6), please note that, following the tradition in Prof. James Ring’s
Infrared Astronomy Department, we credited these papers alphabetically.
Full access to this work is granted for the purposes of study and research.
All other rights reserved.

8This thesis is dedicatedtoHarold May.
This was your dream too, Dad.
.

Figure 0.1 The Zodiacal Light, photographed by the author in Izaña, April 1971.

“Dust in the Wind … all we are is Dust in the Wind …”(Lyric by Kerry Livgren, of Kansas, 1977,
from the album “Point of No Return”)

9Preface and Acknowledgements
This thesis documents a body of spectroscopic observational work on the Zodiacal
Light, which was completed 30 years ago. The results were published in two papers(see Appendices) in Nature, 1972, and in Monthly Notices of the Royal Astronomical
Society, 1974. The project was devised and adopted under the guidance of Prof. James
Ring and Dr. Ken Reay, of the Department of Infrared Astronomy, Imperial College,
London. On my first observing trip I was accompanied by Dr. Reay, and on my second,by a young fellow research student, Tom Hicks. For these publications, we made thedecision to credit our names alphabetically, as “Hicks, May and Reay”, but it isimportant to note here that the conduct of this research was in my hands; it was always
understood that this body of work was my own project, the substance of my PhD.
Alongside this clarification, my respects and thanks are here gladly acknowledged to
my colleagues in this venture. I inherited a Fabry-Perot spectrometer and pulsecounting equipment from Prof. Ring, and spent 18 months entirely rebuilding andupdating both the optics and electronics, in preparation for obtaining essentially the first
viable set of radial velocity measurements, all around the ecliptic, of the Zodiacal Light.
The writing of my thesis was virtually complete in 1974, but the submission wasdeferred due to various pressures. Having returned to the work in 2006, I am happy to
submit a complete and substantially updated thesis in 2007, in which I am now able tooffer many new perspectives, because of the research which has taken place between
1974 and the present day.
Although new observations are currently being undertaken, the core of this thesis is myobservational work conducted between 1971 and 1974. It furthers a line of inquiry
initiated in 1962 by James Ring and others at Mt. Chacaltaya in the Bolivian Andes.
The three years (eventually extended to four) of my initial appointment at Imperial
College Infrared Astronomy Department involved me firstly in an attempt to makeobservations from Testa Grigia (altitude 11,500 ft) in the Italian Alps, totally frustrated
by atrocious weather, and subsequently in the retrieval of the equipment from under six
feet of snow, and its transport back to England. I then concerned myself with theupdating of a spectrometer with its control and read-out equipment, the building of a10small observatory (hut!) to put it in, on a volcanic ridge in Tenerife, and three seasonsof on-site work. The first visit was spent largely in discussing the mysteries of English
blueprints with Tenerife builders, but observations during the remaining two did indeed
provide the first ever comprehensive body of high-resolution spectra of the Zodiacal
Light. Two years of work on the data followed, revealing an unexpected discovery of
Magnesium emission, as well as much new information on the solar Zodiacal Cloud.
I was most fortunate, in the initial phase (1970-4), to work under the wing of thefollowing two experienced Zodiacalists.
My grateful thanks to my initial supervisor, the late Professor James Ring, in the firstplace for taking me on as a research student in the Infrared Astronomy Department ofImperial College, London, subsequently for his constant interest, faith, and advice, andlater for the very great favour, to someone no longer strictly a College member, of
continued access to the life and amenities of the Imperial College Infrared Astronomy
Group in the 1970s.
Also to my co-supervisor Dr. Ken Reay, who gave unsparingly of his time and energy
during the first four years of my studies, at 10, Prince’s Gardens, South Kensington.
My long neglected gratitude to my Dad, Harold May, who strongly supported myscientific endeavours all through my life, until he passed away too young, and too soon
to see the full impact of his guidance and example to me. He also helped with some ofthe diagrams seen here, and his influence, his immaculate attention to symbolic detailand stylish draughtsmanship, can still be seen in some of the graphics in this thesis. I
hope I will at last do him justice. And to my dear Mum, Ruth, who encouraged me in
all aspects of my life, and gave me perhaps the deepest parts of my humanity.
I am deeply indebted to Prof. Francisco Sanchez Martinez of the University of La
Laguna, Tenerife, for the use of the facilities of the Observatorio del Teide in the 1970s,
and for his encouragement and support during my visits to Tenerife as a young student;
to my dear friend, now sadly departed, Dr. Carlos Sanchez Magro, who indispensablyorganised local materials, equipment and labour, and taught me to swear in Spanish! I
hope you will be proud, finally, Carlos, of this fruit of our friendship.

11

To my successor in spectral radial velocity work, Tom Hicks of Imperial College, whoshared part of the observational work, helped in the data reduction, and from whom I
have benefited in many discussions; to J. E. Crabtree, W. Stannard and A. V.
Demenezes for help and indulgence in constructional matters; to Isaac Newton (thedepartmental cat, now sadly departed), my sole comfort and inspiration during theseemingly interminable initial reduction of an enormous pile of data; to the Science
Research Council for financially supporting the first three years research, and to
Stockwell Manor School (where I taught Maths as a supply teacher to fund a further
year of research) and “Queen”, for support and distraction during the final year; to Miss
Christine Mullen (later Mrs. May) for much help and encouragement to sustain thework in circumstances which were sometimes difficult, and to Miss Mary Davies (nowMrs. Mary Lewis), who typed the first thesis manuscript, now lost.
Returning to the study of the Zodiacal Light in 2006, I have based this new version
directly on my 1970s handwritten notes, and in most places it adheres closely to theoriginal text, updated only where qualification was necessary in the light of subsequent
experience. In this second phase, I have benefited from new perspectives and expert
advice of the highest quality, and I offer my sincere thanks:
To Prof. Michael Rowan-Robinson, Head of Astrophysics, Imperial College, London,
who has done me the great honour, as my new supervisor, of enabling me to re-register
for a PhD at Imperial College, offering me the opportunity finally to complete a Great
Circle in my life. Also for giving me, after an absence of 30 years, a ‘home’ in theBlackett building; and above all, for his caring and careful criticism, tutoring andadvice, which have lifted this thesis to Gold Standard.
Once again to Prof. Francisco Sanchez Martinez, a hero to me for 35 years, who has reopened the doors of the Instituto de Astrofisica de Canarias (IAC), to enable me once
again to take advantage of the facilities in the Observatorio del Teide - the place where
this line of research all began - for Charles Piazzi Smyth in 1853, and for me in 1971.
To Dr. Garik Israelian, my prime collaborator in resuming this work, who has paid me
the huge compliment of putting his own research programmes on hold, in order to12plunge into a new Zodiacal Cloud study, and is more than anyone else responsible forhelping me through the final stages of this PhD work. Also to his lovely family.
To Sir Patrick Moore, who inspired me a lifetime ago to be an astronomer in the firstplace, for his personal encouragement to rebuild my participation in Astronomy,
particularly during a three-year period whilst we collaborated on a popular book on the
History of the Universe - Bang! - motivating my return to Zodiacal Light research.
To Dr. Chris Lintott, our third collaborator in the Bang! book, for educating me in theways of modern astronomical communication via the Internet, for a crash course in awhole spectrum of new cosmological concepts, and for many suggestions.
My thanks also, for encouragement and advice in so many areas, to Tom Short, of the
University of Herts., my best friend ever since my old college days, who has been kindenough to check my manuscript; to Prof. Stan Dermott for many invaluable insights; toDr. Tom Kehoe, Ashley Espey, Prof. Ian Robson, Marcus Chown, Robin Rees, Brian
Smallwood, Prof. Tom Gehrels, Yvonne Powell, Carey Lisse, Dr. Noah Petro, and Prof.
Greg Parker; to Jennifer Tunney for truly indispensable proof-reading; to Sara Bricusse
and Claire Bartlett for timetabling and personal backup, and Philip Webb for low-stresstransportation; to Pete Malandrone for technical help, to Justin Shirley-Smith forinvaluable computing support; to Prof. Kirpal Nandra, Prof. Rees Rawlings, Prof. JanetDrew, Juliet Kemp, Loli Sanchez, Henry Jestico, Dr. Elise Laird, Dr. David Clements,
and Diana Moore, all of Imperial College, for smoothing my way back into college life;
And to the examiners of this thesis, Dr. Iwan Williams and Dr. Mark Jones, who, by
their careful analysis, significantly improved my text. To my dear friends Elena
Torrents and Jose Manuel Vergara, and Inge and Patricio Feier, for giving me a spiritualhome in Tenerife; to my children, Jimmy, Louisa and Emily, for their great enthusiasm
for the project, in spite of its all-consuming demands on their Dad’s already stretchedavailability; and to my dear wife, Anita, the essential pivot and inspiration of my life,
without whom I would not have had the will to journey into new lands.
Brian May 2007.
In the reference research for this thesis I have made use of NASA's Astrophysics Data System
Bibliographic Services (ADS). Thanks! I would also like to acknowledge those two new portals to theWorld’s knowledge: Google, and Wikipedia.

13

Figure 0.2 Growth of the observatory in Tenerife over 35 years.
(above) View of the ZL Hut (centre) and the site at Izaña, looking East, 1971.
(below) The Observatorio del Teide, from the same viewpoint, in Sept 2006.

14

List of Figures
0.1 The Zodiacal Light, photographed by the author in Izaña in 1971.
0.2 Views of the Observatorio del Teide in 1971, and 2006.
1.1 Cassini’s paper of 1685, claiming the discovery of the Zodiacal Light.
1.2 A chart from George Jones’ ‘Observations of the Zodiacal Light’ (1856).
1.3 a) The Zodiacal Light seen on March 5, 1856 – after Müller (1865).
1.3 b) The evening Zodiacal Light photographed by Szymanek (2006).
1.4 A simple model of Zodiacal dust density as a function of particle radius anddistance from the Sun.
1.5 Illustrating a line-of-sight integration for prediction of a Doppler-shifted absorption
line spectrum.
1.6 Ingham’s theoretical predictions of a modified Fraunhofer line in sunlight scatteredby a rotating circum-solar dust cloud (1963).
1.7 Low et al’s discovery of the Zodiacal Dust Bands.
1.8 Rowan-Robinson (1991) Grey scale and contour map of 60 micron emission.
1.9 Jones and Rowan-Robinson (1993) Fan and Band model.
1.10 Kelsall et al (1998) Dust cloud model.
1.11 Grogan et al (2001) Dust band model.
1.12 Grogan et al (2001) Simulation of Koronis dust particle ring evolution.
1.13 Reach et al (2003) Ratio of observed ZL spectrum to a black body continuum.
1.14 Reach et al (2003) Comparison of spectra from the ZL and Beta Pictoris.
1.15 Fixsen and Dwek (2002) Overview of the spectrum of the night sky in the entire
range of the infrared.
1.16 Jessberger et al (2001) A porous chondritic IDP collected from the Stratosphere.
1.17 Madsen et al (2005) Spectrum of the Zodiacal Light toward the North ecliptic pole.
2.1 The solar spectrum in the region of the MgI Fraunhofer line.
2.2 Location of the observatory on the island of Tenerife.
2.3 The hut viewed from the South, with coelostat assembly in open position.
2.4 Layout of the ZL building and coelostat.
2.5 a) The coelostat, set on concrete plinth, showing positions of its mirrors.
2.5 b) Showing the path of light through the instrument into the hut.

15

2.6 a) Sky accessible from coelostat marked on a coordinate conversion chart.
2.6 b) Illustrating the stereographic projection.
2.7 Principle of the Fabry-Perot Interferometer.
2.8 Schematic of optics, and control and counting electronics.
2.9 The nitrogen system.
2.10 Coating for maximum transmission.
2.11 The Fabry-Perot in its pressure pot.
2.12 a) The ‘perfect’ Fabry-Perot etalon.
2.12 b) Transmission function of the perfect Fabry-Perot etalon.
2.13 a) The effect of polishing errors.
2.13 b) The effect of coating errors.
2.14 Maxwellian view of F-P plates, and possible cross-section.
2.15 Finesses using a 2” entrance aperture.
2.16 Front end optics with calibration equipment in place.
2.17 Intensity calibration box.
2.18 The New Sampler module.
2.19 Logic schematic of the New Scaler Mk4.
2.20 a) Logic representation of a short period delay.
2.20 b) The equivalent circuit.
2.21 Showing layout of components on Sampler boards.
2.22 Schematic of suitable times to observe the Zodiacal Light in 1971.
3.1 Photometry of intensity calibration.
3.2 Experimental graph of counting rate against voltage across bulb for intensity
calibration set-up.
3.3 The drift method applied to determination of error due to non-uniform illuminationof entrance aperture.
3.4 Determination of the separation of the two calibration lines using the drift method.
3.5 Experimental determination of separation of Zinc and Magnesium lamps.
3.6 Experimental variation of pressure as measured by Bourdon gauge against pressure
controller position in lamp number.
3.7 Illustrating the treatment of data.
3.8 The correction for daylight contamination around dawn.
3.9 Evidence for emission in the data.

16

3.10 Mean night-time emission as a function of time.
3.11 Hourly averages of Mg emission intensity throughout the night.
3.12 Simple model of an airglow layer.
3.13 A ‘difficult’ scan.
3.14 Illustrating Newton’s method of successive approximation.
3.15 Wavelength shift estimates from polynomial fits.
3.16 Comparison of a Gaussian curve with theoretical profiles.
4.0 Example of an unrealistically wide fitted Gaussian curve.
4.1 The Hicks, May and Reay data (HMR) – average wavelength shift, determined by
the method of Gaussian fits, as a function of elongation from the Sun.
4.2 Experimental line widths as a function of elongation.
4.3 Experimental line widths by Madsen et al (2006).
4.4 Experimental line depths as a function of elongation.
4.5 Reay’s predicted curves.
4.6 Predictions of Vanysek and Harwit compared with both sets of data.
4.7 Illustrating the movement of the Sun through the local interstellar medium.
4.8 Showing the two Doppler shifts observed in light arriving at the Earth.
4.9 Showing the geometry of the scattering, in the plane of the Ecliptic only.
4.10 Showing the contribution of an element to the line of sight integration.
4.11 Radial velocity predictions for a linear dust flow model.
4.12 Geometry of Doppler shift of light emission by distant matter.
4.13 Predicted curves for distant emitting matter.
4.14 Comparison of our data with that of Fried (1977).
4.15 Comparison of our data with that of East and Reay (1984).
4.16 Comparison of our data with that of Madsen et al (2006).
4.17 Comparison of East and Reay data (1984) with Madsen et al (2006).
4.18 The result of binning the HMR data in wavelength intervals of 20º.
4.19 Binned folded HMR data, compared with folded WHAM data.
5.1 From Stark (2006) Geometrical representation of the ZC.
5.2 Artist’s impression of how the Zodiacal Cloud was perceived in 1970.
5.3 Artist’s impression of the currently perceived ZC, seen from the orbit of Uranus.
5.4 Showing an impression of an evolved toroidal dust band.

17

Abbreviations and Acronyms
AU......

Astronomical Unit

CIB......

Cosmic Infrared Background

CIBER......

Cosmic Infrared Background ExpeRiment

CTC......

Cis-Terrestrial Cloud

CMB......

Cosmic Microwave Background

COBE......

Cosmic Background Explorer

Dec......

Declination

DIRB......

Diffuse Infrared Background

DIRBE......

Diffuse Infrared Background Experiment

EHT......

Extra High Tension

F-P......

Fabry-Perot (interferometer)

FIRAS......

Far Infrared Absolute Spectrophotometer

GDC......

Geocentric Dust Cloud

HMR......

Hicks, May and Reay

IDP......

Interplanetary Dust Particle

IPD......

Interplanetary Dust

ISM......

(Galactic) InterStellar Medium

IRAS......

Infrared Astronomical SatelliteISO......

Infrared Space Observatory

KBO......

Kuiper Belt Object

LIC......

Local Interstellar Cloud

MDTL......

Motorola Diode Transistor Logic

MNRAS......

Monthly Notices of the Royal Astronomical Society

NPL......

National Physical Laboratory

PTFE......

Polytetrafluoroethylene

RA......

Right Ascension

RHVW......

Rowan-Robinson et al (1990) model of the Zodiacal Cloud

WHAM......

Wisconsin Hydrogen Alpha Mapping Spectrometer

ZDC......

Zodiacal Dust Cloud

ZC......

Zodiacal Cloud

ZL......

Zodiacal Light

18

1.1 The Zodiacal Light – the phenomenon

Chapter 1
Introduction
The scope of this introductory chapter is to review the entire history of Zodiacal
Light research, from the earliest times, through the period of my own original
research, right up to the present day. Particular attention is paid to the developmentof Doppler measurements, and to the most recent revelations of infrared surveys.

1.1

The Zodiacal Light – the phenomenon

The Zodiacal Light, conveniently abbreviated to ‘ZL’, is the name given to a diffuse
cone of light seen in the Western sky in the evening, from around the end of twilight,
and in the Eastern morning sky just before dawn. In the latter incarnation it has beencalled the ‘False Dawn’. If it could be viewed in a completely dark sky when the Sun
was eclipsed, we might imagine that the complete outline of a brightness contour of the
ZL would be an ellipse, with the Sun at its centre, and its long axis roughly aligned withthe ecliptic. To see this phenomenon clearly on Earth, it is necessary to have a good,dark sky, free from clouds, terrestrial light sources, moonlight, or aurorae. It is also anadvantage to be within or near the tropics, since it is here that the cone of light isinclined at a high angle to the horizon, making it still visible when the Sun is well below
the horizon, and the sky is completely dark. The maximum intensity of the ZL isgreater than the brightest parts of the Milky Way, and in ideal conditions, the morningand evening cones are seen to be linked up, by a continuous faint band right around theecliptic. This ‘Zodiacal Band’ is visible all through the night, and shows a secondary
brightening at the anti-solar point, known as the “Gegenschein”. With modern
instruments more sensitive than the human eye, it has become apparent that even welloff the ecliptic, indeed even at the ecliptic pole, there is still some contribution to thesky background brightness from this phenomenon. The Zodiacal Light, in fact,surrounds us.
A knowledge of the nature and distribution of the material responsible for the ZL isinteresting for many reasons. At the time when I began this research, around 1970,
manned space flight was just beginning, and there was a need to evaluate what dangers

19

1.1 The Zodiacal Light – the phenomenon

to astronauts there might be from high-energy micrometeorites. In studies of the originof our Solar System, too, there are benefits in knowing the composition and motions of
the interplanetary dust; the present parameters of this cloud of material can giveindications as to how our planetary system was formed, and how the Sun’s environmentis evolving at present. In the 21st Century, studies of the ZL have become relevant for
new reasons; dozens of exo-planets in other solar systems are now known to us in 2007,
but as early as 1984 in the IRAS experiment, the “Vega Phenomenon” was discovered(see discussion by Beichman 1987) – infrared excesses in radiation indicating clouds ofsolid grains around three nearby Main Sequence stars, Alpha Lyrae (Vega) (Aumann etal 1984), Alpha Piscis Austrinus (Fomalhaut), and Beta Pictoris.

Many otherdiscoveries followed, and it now appears that these exo-circum-stellar dust clouds – theZodiacal aggregations of other suns - are very common features of the environs of stars.
The most studied of the exo-zodiacal clouds is perhaps β Pictoris (Smith, Terile 1984),
whose dust debris disc, viewed conveniently edge-on from Earth, appears to beconcentrated in more than one plane, perhaps indicating disturbance from yet€ thorough knowledge of our own evolving
undetected planets (Hubble website 2006). A
system of planets, debris and dust helps to shed light on the evolution of planetarysystems in general. Of course the search is highlighted by a desire to answer thequestion of whether our Earth, and life as we know it, is unique in Creation, or if suchlife is scattered liberally throughout the Universe. Zodiacal Dust may still be holding
on to some of the secrets relating to these questions.

1.2

Review of observations of the ZL prior to 1970

1.2.1 Earliest views“If I take the wings of the morning, and dwell in the uttermost parts of the sea, even
there shall thy hand lead me, and thy right hand shall hold me” - Psalms 39, vv. 9-10.
This biblical quotation has been mooted (Fechtig et al 2004) as a possible reference tothe appearance of the morning ZL, but other translations give the verse a different
meaning: the ‘Word of God’ translation has it as “If I climb upward on the rays of themorning Sun, or land on the most distant shore of the sea where the Sun sets …”.
Moreover it is common for wings to be poetically ascribed to the Sun and Moon, and20

1.2 Review of observations of the ZL prior to 1970

also to the Morning itself, which may be the case here. So there is some doubt as towhether these words from the Bible’s Book of Psalms actually do indicate an awarenessof the Zodiacal Light in Old Testament times. Are there any other signs of ancientknowledge of the ZL? Fechtig et al (2004) allow that the Aztecs (16th Century) may becorrectly credited with having noticed the ZL, and that there are possible referencesfrom the Chaldeans (600 B.C.) and the ancient Egyptians, but such occasional mentions
are few indeed compared with the liberal scattering of allusions in art of all kinds toother heavenly manifestations (for example, the works of Shakespeare contain a virtual
catalogue of atmospheric phenomena!). The great exception seems to be Omar
Khayyam, the great astronomer-poet of Persia. In his Rubaiyat (circa 1120), there are
two passages highly evocative of the ZL:“Awake! For Morning in the Bowl of NightHas flung the Stone that puts the Stars to Flight”

and …

“When False Dawn streaks the East with cold, gray line
Pour in your cups the pure Blood of the Vine
The Truth, they say, tastes bitter in the mouth,
This is a token that the “Truth” is wine.”
In the West, the scientific ‘discovery’ of the ZL is attributed to Giovanni Domenico
Cassini, the famous French astronomer, who, in his publications of 1683 and 1685 (seeFigure 1.1) claims the find. Certainly he was the first to quantify it, chart it in thecelestial sphere, and, with great insight, suggest an explanation for the phenomenon.
But 22 years earlier, J. Childrey (1661) had already brought the ZL to the attention of
the Western world. In his work, “Britannica Baconica, or the Natural Rarities ofEngland, Scotland, and Whales (sic)” he wrote, “In February you shall see a plainly
discernible way of the twilight striking up towards to the Pleiades … and I believe it …
will be constantly visible at that time of the year. But what the cause of it in Nature
should be I cannot yet imagine but leave it to further enquiry”. This is the classic viewof the evening Zodiacal Light, and to this day, February - March remains the perfecttime of year to see it in the Northern Hemisphere. Cassini and his student Niccolo Fatio
de Duillers went a long way towards ‘explaining’ the appearance of the Zodiacal Light– even suggesting dust particles in orbit around the Sun as the source. But the fact thatthese men were the first to document the extent of this quite spectacular manifestationin the Heavens is in itself curious.

21

1.2 Review of observations of the ZL prior to 1970

Figure 1.1 Cassini’s paper of 1685, claiming the discovery of the Zodiacal Light.

It is remarkable that so few earlier records are available, since, unless a sudden dramaticincrease in brightness occurred in the seventeenth century, the mysterious Light of theZodiac ought to have been very conspicuous in the dark skies of earlier civilisations,comparable in dimensions with the Milky Way, and at its brightest, more luminous.
The number of people living in the modern world who will ever see the Zodiacal Light
must be dwindling fast, as urbanisation spreads with its accompanying light pollution,
but in ancient times the unpolluted night skies familiar to shepherds and sailors ought tohave yielded at least a frequent mention of this ghostly cone of light; and why not awhole folklore, such as pertains to the Milky Way, especially since much of the ancientcivilised world resided in or near the tropics? The very fact that there is no popularlayman’s name in any language for what scientists now call the Zodiacal Light suggeststhat there is an anomaly here. Could it be that our view of the Zodiacal Light is highlyvariable? Cassini was convinced that it disappeared completely between 1665 and
1681, and this man was certainly no casual observer. I have no solution to the puzzle.

22

1.2 Review of observations of the ZL prior to 1970

In the light of Cassini’s reports, along with Jones’s observations detailed below, and the
curious dearth of references pre-Cassini, I am convinced that, despite the lack of any
recent confirmation, we must admit the possibility that the Zodiacal Light has notalways been what it is today. One other early set of observations of the ZL is worthy of
special mention.

Figure 1.2 A chart from “OBSERVATIONS ON THE ZODIACAL LIGHT
from April 2, 1853, to April 22, 1855, made chiefly on board
THE UNITED STATES STEAM-FRIGATE MISSISSIPPI,
During her late cruise in Eastern Seas, and her voyage homeward:
With conclusions from the data thus obtained by George Jones, A.K., Chaplain, United States Navy
UNITED STATES JAPAN EXPEDITION VOLUME III
Washington 1856.”23

1.2 Review of observations of the ZL prior to 1970

On April 2nd 1853, the Rev. George Jones (1856), Chaplain of the United States Navy,set off on a nautical voyage to Japan and back, and made visual observations of theZodiacal Light every night, evening and morning, for the whole duration of the trip, aperiod of over two years. Jones was meticulous in his documentation of the appearanceof the phenomenon, and made efforts to eliminate subjectivity from his findings byseeking corroboration from his companions on the ship, who had no prior expectationsof what they were supposed to see. His records, published in 1856 in Washington, were
the first serious sustained study of the subject, and still represent the biggest single
archive of attentions to the Zodiacal Light; in particular, his accounts of its variability
with time and latitude still give pause for thought. A page of his charts is reproduced in
Figure 1.2, showing Jones’s innovative graphic representation of the sky, his own‘projection’, built around the ecliptic (vertical in the charts), rather than the celestial
equator.
The discovery of the Gegenschein, or counter-glow, is generally attributed to Brorsen,(1854) who first observed it in this year, but he himself acknowledged that Humboldt
had named it in 1803, and it is documented that Esprit Pénéza reported its existence tothe Paris Academy in 1731. It is interesting to note that Jones did not know of theexistence of the Gegenschein until he returned home, but had rejected the observation as
unreliable – because he thought it should not be there - and only then did he realise that
he had seen it one night. This is a reminder that, even in such a careful and unbiased
observer, the tendency is to see what one expects to see, and not see what is notexpected. I will have cause to revisit this idea later. Perhaps we can call it the “Lucky
Find Syndrome” (LFS, for short, of course!). By 1854, early photography was wellunder way, but, although scientists of the day (including Piazzi Smyth) had successfully
secured images of many astronomical subjects, the sensitised plates of the day,
Daguerreotype, Albumen, and Wet Plate Collodion, did not yet possess the sensitivity
needed to record an image of the ZL. Hence we only have artistic impressions to go on
from this period. A fine engraving of its appearance in 1856, published by Müller in1865 (see Fechtig et al 2004), is reproduced here (Figure 1.3a). The Hyades are shownat top left of the picture, and it will be seen that there are two interlopers to the starscene … two planets, right on the line of the ecliptic (I have coloured them orange),
reminding us that the axis of the ZL does indeed lie (roughly) in the ecliptic plane.

24

1.2 Review of observations of the ZL prior to 1970

Figure 1.3a The Zodiacal Light as seen on March 5, 1856,
over the city of Friebourg in Breisgau, stretching up to the Pleiades,
as first described by Cassini; from Müller 1865.

Here the Zodiacal cone ‘strikes up’ towards the Pleiades and Hyades, visible at top left,of the picture, just as in Childrey’s description. For comparison, Figure 1.3b shows aphotographic view, taken about 150 years later. Although the constellation of Taurus isslightly lower in this view, again the Pleiades and Hyades feature, in a strikingly similar
configuration of the Spring ZL, viewed from the Northern Hemisphere. Again, a brightplanet is featured.

25

1.2 Review of observations of the ZL prior to 1970

Figure 1.3b The evening Zodiacal Light, photographed by Szymanek in 2006.
The Pleiades and Hyades are visible, as described by Cassini in 1865, and illustrated
by Müller in1865.

In all ground-based photographs of the ZL that I have seen, atmospheric extinction
modifies the actual recorded shape from a cone to something closer to an ellipse, but
this never appears in drawings; it seems that the naked-eye observer instinctively
compensates for the effect of darkening at the horizon (and reddening, as seen here).

26

1.2 Review of observations of the ZL prior to 1970

1856 was also the year in which the celebrated Scottish astronomer Charles Piazzi
Smyth (1856) published an account of his expedition to Tenerife, to perform site-testingat high altitudes, the intention being to confirm the theory, first proposed by Newton,that the quality of ‘seeing’ would improve at high altitudes. Smyth, one of the firstscientists to apply himself to studies of the ZL, suggested that its spectrum must be
essentially reflected sunlight, and in Tenerife in 1855 he made the first attempts toconfirm this spectroscopically. Over a hundred years later, spectroscopy would be usedsuccessfully to map radial velocity patterns relative to the Earth in the Zodiacal Dust
Cloud, and I’m sure Piazzi would have been happy to know that this too was to takeplace in his beloved Tenerife.

1.2.2 Evidence as to the scattering medium
Anders Jonas Ångström, one of the fathers of spectroscopy, discovered the fundamental
relationship between absorption lines and emission lines, was the first to find evidenceof Hydrogen in the spectrum of the Sun, and gave his name to the Angstrom Unit
(10

−10

m). In 1867 he recorded (incorrectly, it turns out) that he had found the same

blue-green emission line in the spectrum of the Zodiacal Light as he had previously(correctly) found in the Aurora Borealis. Since then, it has been substantially confirmed€that there are no emission lines present in the visible portion of the spectrum of the ZL.
Emission lines from airglow are always liable to be mixed in with ZL measurements, inany ground-based observation, but are now recognisable as such. Medium resolution
spectra of the ‘cones’ (Hoffmeister 1940) showed a continuum with Fraunhofer
absorption lines, and Gindilis (1963) found an energy distribution in the Gegenschein
the same as sunlight, so it may be regarded as established that the ZL shines by sunlightreflected (or more correctly, scattered) by a relatively slow-moving medium. If
electrons were present in the Zodiacal Cloud, due to their high speeds, they would add ahighly ‘Doppler-smeared’ component to the scattered spectrum, filling in the absorption
lines. This effect was shown to be absent by Blackwell and Ingham (1961), who, bymeasuring the depth of Fraunhofer lines, placed an upper limit of 120/ cm 3 on electron
density at 1 AU, so the ZL material must be almost exclusively dust. In the light oflater ZC models, perhaps this conclusion is now in doubt, since models with orbits that€are elliptical, rather than circular, also show shallower absorption lines, and some27

1.2 Review of observations of the ZL prior to 1970

observers have now claimed to have confirmed this observationally (Madsen et al 2006
– see Chapter 4). The composition of the dust was variously estimated to be ‘irons’ and‘stones’ (predominantly silicates), similar to meteorites found on Earth, or graphites, (assome exo-planets are now believed to be), or a mixture of ices (solid carbon dioxide,ammonia, water, etc.), though these would have short lifetimes within a few
Astronomical Units from the Sun. The shapes of particles actually collected were, andare, predominantly irregular, or ‘fluffy’, though a few approximate better to the smooth
spheres almost invariably used (Hemenway 1967) in early theoretical models!

1.2.3 Measurable quantities in ZL observations
Measurements of the Zodiacal Light available to various observers up to the 1970’s may
be summarised according to type as follows:a) Photometry - Measurements are made of the intensity of the ZL as a function ofangular distance (‘elongation’) from the Sun, and distance from the axis, across the‘cone’. This information can provide pointers to the physical parameters of the cloud,discriminating between theoretical models of particle size and distribution (see, forexample, Smith et al 1965). Measurements of colour, again as a function of elongationfrom the Sun, have provided clues to particle sizes (e.g. Elvey 1934). Notably, Carlos
Sanchez Magro of the University of La Laguna in Tenerife (PhD Thesis 1974), in aseven-year study, found a dependence of Zodiacal Light brightness on solar activity.
b) Polarisation - Measurements are made of the polarisation of the ZL, usually as afunction of elongation from the Sun, an indicator of particle sizes and properties (e.g.
Sanchez 1969).
c) Solar Corona measurements - The outer regions of the Solar Corona, so stunningly
visible at times of total solar eclipse, merge into the inner parts of the ZL, so
measurements on the Corona can be relevant (Beavers et al 1980).
d) Position of the axis of the ZL - measurements are made of the angle the axis makeswith the ecliptic and its points of intersection. The aim is to ascertain which plane the28

1.2 Review of observations of the ZL prior to 1970

dust principally lies in, and if its distribution is symmetrical about this plane, andsymmetrical East to West. A dependence of the axis position on the latitude of theobserver would indicate a GDC component (Geocentric Dust Cloud) – in other words,
that a significant proportion of the dust was located near the Earth, rather than in orbits
around the Sun. Thus these measurements can enable us to inform an opinion about
where the dust lies (e.g. Huruhata, Masaaki 1965). Wolstencroft et al (1967) concluded
that no GDC component had been detected.
e) Absorption line strength measurement - the spectrum of the Zodiacal Light (sunlightreflected from moving particles) is compared with that of the Sun. The lessening inspectral line strength was believed to indicate a proportion of particles with highvelocities ‘smearing ‘ the scattered spectrum (Blackwell and Ingham 1961).
f) Space-borne detection and collection of particles – samples are collected from spacevehicles, from the Moon, from deep sea sediments, from permafrost in Antarctica, andfrom the Stratosphere (see section 1.4.3).
g) Emission (black body) spectra - the warm Zodiacal dust emits radiation in theinfrared. This can be detected by ground-based observations, but only through theavailable atmospheric ‘windows’. But this line of enquiry is much easier from avantage point in space, where there is no restriction on the wavelengths which can bestudied. This kind of observation came of age much later, with the IRAS satellite andothers (see 1.4.2).
h) Radial velocity measurements - the speed and direction of the motions of dust are
inferred from Doppler shifts in the absorption lines in the scattered solar spectrum. Thisis the line of inquiry pursued in this research project, again principally as a function ofelongation from the Sun. Ground-based spectroscopy of the Zodiacal Light, such asthat described in this thesis, is well-known to be difficult: there are appreciable
challenges to meet – among them, extremely low light levels, atmospheric extinction,
airglow, and light pollution from solar and lunar twilight. It would clearly be very
attractive to mount a corresponding experiment in space, where these problems would
not occur; however, on the ground, experiments can be conducted over a long time
period, and much data accumulated.

29

1.2 Review of observations of the ZL prior to 1970i) Shapes of Spectral Lines - Clarke et al (1995) computed actual profiles due tointegrations of Doppler shifts, and were able to predict a bifurcation of minima at
elongations close to the Sun. As it becomes possible to observe at higher spectral
resolutions, much more could be done in this area.

1.2.4 Physical parameters and models of the ZL
The variables of interest in the interplanetary medium are the number density, size,
shape and albedo of the particles, their distance from the Sun, and their speed anddirection of motion, almost all of which may be a function of any of the others. In
addition, we would like to know the chemical composition of the material, what
proportion is gas and what dust, what are its rates of evaporation, melting points, anddegree of ionisation and so on. The reason for the slow progress in early ZL research isimmediately apparent, since any model constructed to fit a comparatively small range ofpossible experimental measurements is required to yield values of an enormous numberof interconnected physical variables. Simple models of the Zodiacal Cloud are basedon number density as a function of position in the Solar System and particle radius.
Some early observers attempted to attribute the existence of the ZL to scattering of lightin the Earth’s atmosphere, or (as favoured by Jones in 1865) from a dust cloud centredon the Earth; but even in 1683 Cassini had suggested that its appearance might be due tosunlight scattered by a lens-shaped (lenticular) circumsolar cloud of dust concentrated
in the plane of the planets, the Earth being situated some way out towards the edge ofthe ‘lens’. This, aside from some exciting recent discoveries about the fine structure of
the cloud – zones, or bands, of concentration of material attributed to asteroidal erosion
and planetary gravitation - is roughly the view of most astronomers today, though there
is still a lingering doubt as to whether a component might be attributed to a Geocentric
Dust Cloud (GDC). The Geocentric Theory, though not popular by 1970, still had its
adherents (Peale 1966; Divari 1965 and 1966). Early satellite measurements of particle
densities seemed to give support to the existence of a significant dust cloud extending toa few Earth radii, giving cause for alarm to those contemplating manned exploration ofSpace. However, these were later shown to be suspect by Nilsson and Southworth(1967). These researchers, in 1972, detected only TWO impacts in their detector tubes

30

1.2 Review of observations of the ZL prior to 1970

aboard the OGO 2 satellite in hundreds of hours, and the reality of even these was indoubt! So they were able to put an upper limit on micrometeoroids in the vicinity of the
Earth of mass greater than 10 −12 g, of 0.01 particles per cubic metre per second (1967).
Indeed these early reports of high concentrations of dust in the region of the Earth were
never since confirmed, though we shall see later (1.2.4) that there now seems to be€€evidence for at least a circumsolar band of dust in the region of the Earth’s orbit.
Searches for parallax effects in the ZL axis dependent on the observer’s latitude yielded
conflicting results (perhaps hardly surprising given the difficulty of correcting foratmospheric extinction, and the extreme difficulty of seeing the ZL at all in highlatitudes). Roosen (1970) failed to detect the Earth’s shadow passing through theGegenschein, an effect which ought to have been detectable if the dust were near to the
Earth (except on one occasion, when a comet was held to be responsible). Roosen alsoestablished that the predicted L5 cloud, at the Earth-Moon Libration point, was nonexistent. The exhaustive theoretical treatment of Shapiro, Lautman and Columbo (1967)
concluded that the existence of a significant GDC is ‘unlikely’.
Most researchers now favour circum-solar models, and these have frequently used ZL
photometry and polarimetry measurements as their starting points (Behr and Seidentopf
1953; Divari and Asaad 1959; Ingham 1961). Ingham included solar corona data, andfitted parameters to a compromise curve, since there was some difficulty in bridging thetwo sets of data. The general method, in fitting models to experimental findings, is tomake certain simplifying assumptions about the particles (typically spherical particles incircular orbits, with uniform specific gravity and albedo), to restrict attention to theplane of the ecliptic, and to assume values of some parameters, in order to make asuitable fit to the observations by adjustment of the others.
A representative model is one in which the albedo (ratio of reflected light to incident
light) of the particles:

γ = 0.2 – 0.5
and the dependence of number density ( n ) on particle radius ( a ) and distance from the€Sun ( r ) has the form:

€

€

€

31

1.2 Review of observations of the ZL prior to 1970

n(a,r ) = n0 a − p r − α……

equation (1.1)

where p = 3, 4, 5, and α = 1, 3/2, 2.

€This is already, evidently, a simplified model, since in reality there is no reason to€€ p and α are constant, or even orthogonal with each other. The value of
suppose thatno is chosen so as to give a total number density of 240 particles per cubic km at 1 AU
from the Sun, and a lower limit of 0.3 microns is imposed on the particle radius. This is€€because very small particles are ejected from the Solar System by the pressure of€sunlight (Giese 1961); the value of amin can in fact be 0.2 microns to 0.4, depending on3

the density of the particle, which is usually estimated as between 1 and 3 g cm .
Further, particles less than 0.2 microns in radius would scatter light selectively€

according to wavelength, giving rise to predictions of colourations in the ZL which havenot been observed, though, again, observations are not all in €agreement. In thisanalysis, the particles are large compared with the wavelength of light under discussion,
so the complex theory of Mie Scattering is not needed. Values of p between 0 and 5
were favoured by Weinberg (1965) and also by Ingham (1961), from photometry andpolarimetry measurements (see Figure 1.4).

€
Turning attention to the dependence of n on distance from the Sun, Over (1958)
predicted a dust-free zone within about 4 solar radii (4 rs ) owing to vaporisation by theSun. Infrared emission peaks were reported by Petersen (1967) supporting the€prediction. This can be incorporated in the model most conveniently by specifying two€4 solar radii) in which n(a,r ) varies as
domains for relation 1.1, an inner one ( r <(4rs − r)−α(Figure 1.4).

€€

where α is large, and the outer region where the model is as before

€

€

32

1.2 Review of observations of the ZL prior to 1970

Figure 1.4 A simple model of Zodiacal dust density as a function of particle
radius and distance from the Sun.

Most in situ ZL measurements in this period were made from satellites in theneighbourhood of the Earth, and served to verify and extend the Earth-bound data.
Hanner and Weinberg (1973) conducted an interesting experiment in whichmeasurements were made of the Gegenschein intensity as a function of distance fromthe Sun from an outgoing space vehicle. Their results lent support to the model:n( r ) ∝ 1 r….

i.e. equation 1.1 with α =1.

The reason why the Gegenschein is there at all was certainly not clear prior to 1970,€€ and may still be in question. If it is not due to a GDC, or to Lunar ejecta streaming
outwards in the solar wind (Jeffries 1971), perhaps n(r ) in the circumsolar cloudincreases outside the Earth’s orbit, or perhaps the effect is caused by a back-scatteringpeak in intensity which might be expected from ‘fluffy’ particles, and has been€
predicted from models using dielectric materials
in general. In none of the above

33

1.2 Review of observations of the ZL prior to 1970

scenarios is the observed shape of the localised maximum at the antisolar point easy to
explain.
Before leaving this sketch of the state ZL knowledge in 1970-4, when my experiment
took place, it is interesting to note that the era of space-borne observations of theinterplanetary medium was just beginning, and at this stage, things were not going aswell as might have been hoped! Sparrow and Ney (1973), based on data from OSO-5,
concluded that there was essentially no variation in the brightness or the polarisation of
the ZL from 1969 to 1973. Levasseur and Blamont (1973), on the other hand, using
data from satellite OS2 “Tournasol” disagreed entirely. It might be interesting (butoutside the scope of this work) to revisit these two papers in the light of what we now
know from the 1984 IRAS revelations. This, coupled with the similarly inconclusive
results of particle collection mentioned earlier, set a tone of uncertainty in the results ofin situ studies, which around 1970 gave good reason to persist in Earth-based researchon the ZL, in particular with the relatively new technique of Doppler radial velocity
measurements.

1.2.5 Doppler investigations of the ZL prior to 1970
The approach of radial velocity measurement is especially valuable in the study of theZL, since it provides completely independent estimates of some of the parameters – viz.
speed and direction distributions in the cloud, leading to insights into its formation andevolution.

Christian Doppler, an Austrian mathematician and physicist working in Prague, was thefirst to draw attention to the change of the pitch of sound waves due to a relative motionof the emitter and the observer, an effect very noticeable in the siren noise on a passing
police car, and thereafter known as the Doppler effect. He was able to offer an
explanation, in terms of the way the wavelengths perceived by the observer were‘squeezed or stretched’ by the relative motion, and proposed, in his paper in 1842,
translated as “On the Coloured Light of Double Stars and Certain Other Stars of the
Heavens”, that a similar effect would be seen with light waves. It is easy to visualisethat light waves emitted from a moving source are squeezed closer together (as34

1.2 Review of observations of the ZL prior to 1970

perceived by us) if the source is approaching us, producing a shift in wavelength to theblue, and stretched out if the source is moving away, so that we see a red shift. In 1848,
French astronomer Hippolyte Fizeau, the first to make a serious estimate of the speed oflight, confirmed experimentally that the Doppler effect changed the positions of spectral
lines, by amounts proportional to a source’s radial velocity relative to us.
The Doppler shift Δλ from a source of light of wavelength λ , moving at a radial
velocity V relative to us, is given by the simple relation

€

λ
V
Δ€
=
λ
c

€
where c is the speed of light.

In 1862, Huggins and Miller remarked that the analysis of Fraunhofer lines in stars€€

could lead, not only to analysis of their material compositions, but also to estimates of
their motions. “If the stars were moving towards or from the Earth, their motion,
compounded with the Earth’s motion, would alter to an observer on the Earth therefrangibility of the light emitted by them, and consequently the lines of terrestrial
substances would no longer coincide in position in the spectrum with the dark lines
produced by the absorption of the vapours of the same substances existing in the stars.”Thus was born the science of spectral radial velocity measurements.
We can note in passing that the Doppler effect became a pillar of the evidence for the‘Expanding Universe’ as we perceive it today. The above relationship was invoked by
Vesto Slipher in 1917, who found that the light from the 21 of the 25 galaxies hestudied showed a red shift, indicating that they were receding from us, and in 1929
Edwin Hubble was able to show that the speed of their recession was proportional totheir distance from us.
Much closer to home, we have seen that Blackwell and Ingham (1961) used the conceptof ‘smearing’ of absorption lines due to high Doppler shifts to place an upper limit on
free electrons in the interplanetary medium. This introduced Doppler techniques to thestudy of the Zodiacal Light, and it was Ingham who laid the foundations for use of the
Doppler method for investigations of the motions of interplanetary dust.

35

1.2 Review of observations of the ZL prior to 1970

The profile of an absorption feature in the ZL spectrum in a particular area of sky
corresponds to a unique line of sight through the Zodiacal Cloud, and is a combinationof the effects of all the light-scattering dust particles along this line. We can regard thespectrum as the summation of contributions from all the elements, or ‘cells’ in the lineof sight; each has a radial velocity relative to the Earth which is a combination of its
own motion and that of the Earth, and leaves its ‘signature’ on the resulting combined
spectrum. By comparing such a theoretical prediction with experiment, for various
different directions relative to the Sun, information can be derived on relative number
and size distributions and velocity distribution in the cloud, quite independently of otherobservations. Typical geometry applying to a radial velocity element, or ‘cell’, isshown in Figure 1.5.

Figure 1.5 Illustrating a line-of-sight integration for prediction of a Doppler-shifted absorption line
spectrum from a rotating dust cloud model.

36

1.2 Review of observations of the ZL prior to 1970

This view from ‘above’ the plane of the ecliptic shows the geometry for a cell of dust,
viewed from the Earth, along a line of sight at an elongation angle ε from the Sun. In
the particular case illustrated, the dust is moving with speed Vc , in a circular solar orbit,
and therefore has no radial velocity relative to the Sun (in fact the orbits may not be€

circular, but may be nearly so). The cell ‘sees’ the continuous visible light emission€ lines of absorption, with zero shift.
spectrum of the Sun, crossed by dark Fraunhofer
Dust in the cell scatters light towards the Earth. But this cell is moving away from the
Earth radially with a component of its velocity along the line of sight Vc sin θ , while
the Earth’s orbital motion, velocity VE , produces an opposing, closing component ofvelocity along the line of sight VE sin ε . The resulting radial velocity to be used in€applying the Doppler shift is thus Vc sin θ - VE sin ε . We can eliminate the cell’s€velocity as follows. We know that for circular orbits:€
Vc 2 R= ,VE 2 r€€where r is the distance of the dust cell from the Sun, and R is the distance of the Earth€ the Sun. In the Earth-Sun-Cell triangle the Sine Law gives:
from

R
r
=
, or
sin ε sin θ

€

hence
€

R sinθ €=r sin ε

Vc
sin θ
= €
VE
sin ε

And the required relative radial velocity ΔV is given by€
ΔV = VEsinθ
€ε sin θ − VE sin ε .
sin

The resulting contribution of this cell to the spectrum of the ZL is light with a red shift€λ , where:
of Δ

€

37

1.2 Review of observations of the ZL prior to 1970

λ  sin θsinθ − sin ε Δλ = VE c  sin ε
Ingham first arrived at this formula in his paper (1969) in a slightly different form,€ using cosecs.
An integral of this expression for all cells along the line of sight, weighted according to
the predicted density of dust for each cell position, and the percentage of incident light
re-directed towards the observer, gives an ‘average’ radial velocity, and hence anaverage wavelength shift for any absorption line, which can be compared with the shifts
found by experiment. Clearly this is an approximation; a full treatment integrates theintensity at each wavelength, and predicts the profile of the absorption line in detail.
In 1963, Ingham (1963) published his theoretical predictions, calculated in this way, ofthe manner in which the shape of a solar Fraunhofer line would be modified byscattering from a rotating dust cloud medium (Figure 1.6), in the ecliptic plane, as afunction of elongation from the Sun. He showed that a broadening and overall shift of
the wavelength of the line ought to be detectable, giving information with which to map
the motions of Zodiacal dust. All that was now needed was the observational evidence
to compare with this theory.

Figure 1.6 Ingham’s (1963) theoretical predictions for the shape of a modified Fraunhofer line
(elongation 30º) in sunlight scattered by a rotating circum-solar dust cloud.

38

1.2 Review of observations of the ZL prior to 1970

Attempts to record a high-resolution spectrum of the ZL as a tool for revealing anddiscriminating between motional signatures, are limited only by the resolutionluminosity (R-L) product of the instrument and time available to spend on observing
each point. The need for high R-L product was a good reason for choosing aninterferometric spectrometer, the Fabry-Perot etalon (Jacquinot 1954), but as we willsee, there are other attractive alternatives.
Beginning in June-August 1962, under the leadership of Professor James Ring at
Chacaltaya in Bolivia, Hindle, Reay and Ring (1968) employed a Fabry-Perot etalon inseries with an Ebert Spectrometer, with a resulting resolved width of about 1.2
Angstroms (Å), to record Zodiacal Light spectra in the region of the H β absorption line
over a range of about 10 Å, with the intention of determining Doppler shifts in this line.
Detection was by cooled photomultiplier and counting electronics, a precedent followed€ the absorption line they had
in many subsequent experiments of this type. Instead of
hoped to see, they found strong evidence of H β emission. Only a small number ofobservations were made, mostly of the evening ZL, and the conclusions were:€ H line in the spectrum of the Zodiacal Light did(a) That the profile of the Balmerβ
not match that of the Sun, either in position or shape.
(b) That some scans showed clear evidence of emission at the H β wavelength,€

especially near the Galactic plane, and it was proposed that this emission wasfrom Hydrogen in the Galactic plane.

€
In 1964, Prof. Ring’s team made a new set of observations, using a 3” Fabry-Perot andan interference filter premonochromator (Daehler et al 1968). The width of the passband of the system was 1.5 Å, and about 50 scans were obtained of morning andevening ZL. The most immediate result was the confirmation of H β emission, whichstill obscured the absorption line. The emission was stronger in the morning than in theevening, and apparently had a maximum intensity at 70 degrees elongation from the
Sun in the ZL. So strong was the emission that it was€difficult for these observers to
draw any conclusions about the obscured absorption line shape.

39

1.2 Review of observations of the ZL prior to 1970

In November and December 1967, Reay and Ring repeated their observations from asite at Testa Grigia in the Italian Alps, looking at regions of the sky clear of the ecliptic,
in both morning and evening, with an improved instrument. They found more H βemission, and not only in the plane of the Milky Way. An automatic stepping systemand careful calibration were introduced so the identification of the H β wavelength was€

now beyond doubt. Even in a moonlit sky (contributing, of course, reflected sunlightwith an unshifted H β line in absorption) the emission was strongly seen. Shifts of the€line were observed in some scans in the evening Zodiacal cone – all towards the blue.
This was still far from the simple picture of a shifted absorption line which the€experimenters
had looked for.
This time, however, Reay and Ring attempted to subtract the emission lines from thespectra, to reveal an interpolated absorption feature which hinted at the shifts they werelooking for. These correlated tolerably well with a prograde dust model, but theprobable errors were high, and the data was confined to evening ZL elongations
between 40 and 120 degrees. Nevertheless these data (plotted as green circles in Figure4.1) stood for some years as the results with which motional models were compared. In
their letter to Nature in 1968, Reay and Ring announced their intention to overcome theproblems of obscuration of the absorption feature by H β emission, by switching theirattentions to the MgI line at 5183.6 Å.

€
Returning to the absorption line data, Bandermann
and Wolstencroft (1969) attempted abetter fit to the Reay and Ring data, by basing their model on elliptical dust orbits,
rather than circular, but the differences in the predictions were small, and Reay (1969)
produced, theoretically, larger variations by including the effects of radiation pressure
on orbital speeds. These effects are large in the case of sub-micron particles. Other
attempts were made to construct the theoretical shape of an absorption feature fromrotating dust-cloud models, notably James (1969).
Vanysek and Harwit (1970) pointed out that the Reay and Ring data were actually notinconsistent with the theory that a ‘Cis-Terrestrial Cloud’ (CTC) was responsible for the
ZL, and predicted a drop in the wavelength shift at elongations of about 70 degrees.
They also showed that the application of full Mie Theory to the scattering problem

40

1.2 Review of observations of the ZL prior to 1970

produces predictions very similar to the Hulst-Ingham-Reay method of treatingreflection and refraction separately (Ingham 1961).
New experimental data was obtained by James and Smeethe (1970), using the newtarget absorption line, MgI (5183.6 Å), who attempted to make observations of the moreelusive (because of twilight and atmospheric extinction) lower elongation regions in theboth the morning and evening ZL. Their results (shown as brown crosses in Figure 4.1)
did show high shifts at low elongations, which is what the circumsolar prograde models
predict, but the radial velocities they found looked too large to lie on the p = 5
theoretical curve (see Chapter 4), and again the probable errors made it impossible tosupport any firm conclusions. The observations were limited to a very small range of€elongations, so no significant curve-fitting could be done to discriminate
betweenmodels. James and Smeethe concluded that their observations were consistent with adust distribution model in which p < 3 , and most of the dust was in prograde solar
orbits, but the matter remained unproven.

€ prior to 1970, it had been hoped that the radial velocity
Throughout these studies
patterns found would discriminate between dust in prograde and retrograde orbits
around the Sun: if the dust were determined to be moving in orbits around the Sun in
the same sense as the Earth (‘prograde’), an asteroidal origin would be favoured. In thishypothesis, due to Fesenkov (1958) collisions and/or cosmic ray bombardment within
the asteroid belt produce debris, which slowly spirals in towards the Sun. Other authorsof the period, notably Whipple, favoured cometary debris as the principal component of
Zodiacal dust, an idea first suggested by Fesenkov in 1914. Since many comets (in fact
perhaps a majority of those with long periods) have retrograde orbits, presumably thediscovery of any retrograde elements must favour a cometary component. It isinteresting that these two opposing views of the origins of interplanetary dust are still
hotly defended in 2007 (see Chapter 4). A Zodiacal Cloud origin coeval with the Earthand other planets was thought by most to be unlikely in 1974, since the gradual collapseof the orbits of small particles, and their probable eventual vaporisation by the Sun(Dobrovolskii and Egibov 1973), over a timescale short relative to the age of the Solar
System, requires that the material is continuously replenished. If the source ofreplenishment of the material is outside the Solar System, we might expect an entirely

41

1.2 Review of observations of the ZL prior to 1970

different pattern of velocities. However, opinions on the relative strengths of thevarious forces acting on a particle vary greatly, so the question of the age of theZodiacal dust may be regarded as still unresolved, and certainly was in 1970. If the dust
did turn out to be a remnant of the creation of the Sun’s whole planetary system, we
would still expect a predominantly prograde cloud.
To sum up, in 1970, uncertainties in observations of radial velocities, combined withuncertainties as to which theoretical model should be used as a comparison, left thequestion of the real motions of Zodiacal Dust completely open.

1.3

Review of the period of my observations of the ZL,
1970-1974

It was the intention of this PhD study to fill in the gaps and uncertainties left inobservational data in 1970; to make a comprehensive survey of a Doppler-modified
Fraunhofer line in the ecliptic plane at all elongations from the Sun, and hence constrain
an accurate model of the Zodiacal Cloud’s configuration and velocity structure. The
spectral line chosen was the MgI 5183.6Å line, because of its (hoped) freedom fromemission line contamination, and a new pressure-scanned Fabry-Perot interferometer
was designed and built to make the observations.
In 1972 and 1974, we, Hicks, May and Reay of Imperial College Infrared Astronomy
Department, under the supervision of Prof. James Ring, published, in two papers, theresults from the research which forms the basis of this thesis. The first paper in Nature(Hicks, May, Reay 1972) reported the MgI emission which had unexpectedly turned up
in the new spectra. The second paper in MNRAS (Hicks, May, Reay 1974) reported theevidence, found in our Zodiacal Light spectra, of radial velocity shifts in the MgIabsorption line, the first comprehensive mapping of such shifts ever achieved. The
paper confirmed for the first time with any certainty that the pattern of Doppler shifts
was consistent with a high proportion of circum-solar dust moving in prograde orbits,
though some significant asymmetries between morning and evening velocities werenoted. A third paper detailing further interpretation with reference to dust cloud42

1.3 Review of the period of my observations of the ZL, 1970-1974

models, including a putative interstellar component, was never completed. The dataobtained from this project, along with some comparisons with simple models of the
ZDC, are detailed in Chapter 4, which constitutes the first full publication of the resultsof my research 1970-4.

1.4

Review of recent developments in ZL research,
1974 – 2007

1.4.1 Overview
Returning to the field of Zodiacal Light research after 30 years, I found I was in a good
position to obtain a perspective on what progress had been made in the intervening time.
I am indebted to four excellent papers for much of the information in this review –Grogan et al (2001), Reach et al (2003), Kelsall et al (1998), and Sykes et al (2004).
My first observation was that, in keeping with the acronymic style of pigeon-holingnow so universal, in which everything from espionage organisations, to diseases, tofurniture warehouses, find themselves summed up in three letters, interplanetary dust
particles are now referred to as “IDP”s. And the business of detecting them, and theircomposition, and orbital elements, has progressed enormously; we even have some inour hands, to study on Earth.
Secondly, it soon became apparent that there are now historically two distinct eras
within modern ZL research: pre-IRAS and post-IRAS (InfraRed Astronomy Satellite).
In the period 1974-1982, observations were extended, and the existing dynamical
models updated, but the overall picture remained essentially unchanged, except forsignificant advances in in-situ collection techniques, and the analysis of meteoric
impacts on Lunar samples (Singer 1969). It was the data gathered from the IRAS
satellite in 1983 (while my rock group Queen was touring football stadia in SouthAmerica) which finally moved our picture of the Zodiacal Cloud on, in a giant step,from diffuse oblate spheroid, the limit of anyone’s imagination since the 17th century,
to a structured, dynamic and complex dust ensemble.

43

1.4 Recent developments in ZL research

Figure 1.7 Low et al’s discovery of the Zodiacal Dust Bands. For details, see Low et al (1984).

This features the now well-known ‘IRAS Bands’ (Low et al 1984), (see Figure 1.7), of
which at least 4 pairs are now known, contributing about 10 per cent of the total
Zodiacal emission in the infrared, and, on a small scale, cometary dust trails (around 8of these have been seen at the time of writing). Moreover, it revealed a local dust cloudof a kind near the Earth, but the idea of a GDC was, at a stroke, replaced by a local
toroidal dust intensification near the Earth’s in its path, due to resonance in the dust
orbits; effectively a ‘dust lane’ stretching out behind the Earth in its journey around theSun. This research also revealed (Hauser et al 1984) that the axis, the line of brightest
scattering of the ZL, changes its orientation to the ecliptic plane, depending on where
the Earth is in its orbit. The position of the axis varies in a yearly sinusoidal cycle witha maximum tilt of about two degrees, bringing about an expected seasonal intensity
variation of around 15 per cent, in addition to a 5 per cent variation in a different phase,resulting from the offset of the Sun from the centre of the cloud. This, strangely, brings
to mind once more the findings of the patient Rev. Jones in 1856, who did see avariation over the two years in which he made his naked-eye observations of thecontours of the ZL as he saw them from his boat. He saw what he thought were44

1.4 Recent developments in ZL researchparallax-induced variations in the position of the axis, but since his voyage moved him
not only in latitude, but around the Sun too, over a period of two years, it is perhaps
possible that his maps record the first documentation of the inherent asymmetry of the
ZC with respect to the ecliptic plane.
Dermott et al (1992) achieved an impressive theoretical fit to the yearly cyclic variation,using their own updated ‘SIMUL’ model of asteroidal dust of radius 3.4 µm , spirallingslowly inwards under the influence of Poynting-Robertson drag (Robertson 1937), andperturbation of the planets. A little later, Leisawitz et al (1994) using the new data from€ satellite, extended the
IRAS, but also that from the DIRBE survey aboard the COBE
analysis to link the dust-band data with particles of varying sizes produced by collisionsof asteroids in the Eos family. All the new analyses led to a much more detailed
picture: a finely tuned version of the ‘traditional’ lenticular-shaped ZC, now
symmetrically distributed not about the ecliptic, but relative to a plane dictated byplanetary perturbation, inclined at about 2 degrees to the ecliptic, and centred on a point
0.13 AU from the centre of the Sun. This updated view of the ‘traditional’ ZC wastheoretically modelled with a combination of cometary and asteroidal contributions, theproportions of which are still under debate. But the new picture also included a threedimensional notional map of the newly discovered, relatively stable ‘bands’ (attributed
to toroidal concentrations of material), oriented at roughly one degree angles to theplane of the lenticular cloud, identified by Low et al (1984) and studied in detail by
Grogan et al (2001) with matter spiralling in, under the influence of PoyntingRobertson drag, from collisions in certain families of asteroids. Added to this are thediscrete trails of dust left by comets, seen unmistakeably for the first time in the IRAS
composite images, (Sykes et al 1986) and a circumsolar ring at just outside 1 AU, aspredicted by Gold, in 1975, observationally supported by Dermott et al (1994), andconfirmed by Reach et al in 1995. The IRAS observations showed a consistent 3 percent increase in intensity at all its wavelengths behind the Earth, looking directly at its
line of motion, relative to that seen in front of the Earth in its orbit. This is consistent
with the predictions for a resonant ring of width 0.4 AU with a cavity in front of the
Earth. Dermott et al noted that this must produce an increased rate of capture in the
Earth’s upper atmosphere, and that a seasonal variation ought to be seen, according towhere the Earth lies in relation to the Zodiacal Cloud. They also saw this as a possible

45

1.4 Recent developments in ZL research

connection to the accretion of Life-bearing carbonaceous particles in the history of the
Earth. Those particles not captured by the planets, of course, suffer a different fate:under the influence of Poynting-Robertson drag, and solar wind braking effects, theZodiacal Cloud loses one ton of material per second, as the dust spirals ever inwards, tobe vaporised at a distance of a few solar radii from the Sun.
The researches of these teams, as noted by Kelsall et al (1998), were all initiallydirected towards evaluations of infrared background emissions, from Galactic andextragalactic sources. It was a need to eliminate the foreground ‘contamination’ by the
local interplanetary dust which led to their careful analysis of the ZC. Rowan-Robinson
et al (1990) derived, from early IRAS data, a model of 12-100 micrometre emission by
the Zodiacal Cloud, which I will refer to hereinafter as the RHVW model, which they
were later able to use to subtract from the observational maps, to show the ‘shoulders’of the IRAS bands (see details in 1.4.2.1).
So the Zodiacal Light can certainly no longer be thought of as smoothly symmetricalabout the ecliptic plane. IRAS’s data collection was in the infrared, at 12, 25, 60 and
100 µm , and all these observations are of emission, rather than scattering of sunlight,
but there are now claims, from Ishiguru et al (1999), that corresponding features due tothe existence of the IRAS bands exist in scattered light at visible wavelengths. They€scanned the surface brightness of the ZL at right angles to the ecliptic plane, anddiscerned the ‘shoulders’ of discrete bands, superposed on the smooth curve ofbrightness associated with the ‘traditional’ lenticular ZC. This research begs theinteresting question of why 50 years of careful photometry of the ZL from the Earthnever revealed these ‘shoulders’ before. Perhaps this is another example of LFS – theLucky Find Syndrome – nobody saw the ‘shoulders’ because nobody thought theyshould be there. Ishiguru et al are also one of the few modern observing teams tofollow up on the Gegenschein, and claim to have detected parallax in the optical bands
in the Gegenschein area, placing the source material at 1.6 AU from the Sun. In a
recent collection of accounts of progress in ZL observations, the book “InterplanetaryDust” (Grün et al 2001), the Gegenschein receives only four mentions, three in an(excellent) historical review, and one in the glossary.

46

1.4 Recent developments in ZL research

At this point, since the newly discovered detailed structures of the ZL must in future
play an essential role in constructing models of radial velocity patterns, it is useful tolook at the important recent advances in more detail.

1.4.2

Infrared Surveys

1.4.2.1 Modelling based on IRAS data
Rowan-Robinson et al (1990) investigated the distribution of Zodiacal dust as afunction of position out of the plane of symmetry. Following Leinert (1985) and others,
they looked for a fit to the IRAS flux measurements, using a functional dependence of
number density of the dust grains: r − γn(r ) = n0   f (β0 ) , A

€where n0 is the number density near the Earth’s orbit, A = 1.AU, β0 is the elevationfrom the symmetry plane of the dust, and γ ≈ 1.3 from Helios data (Leinert et al 1981).
A popular form of the function f (β0 ) was the ‘Fan’ model:€€€
f (β0 ) = exp(−α sin β0 ) ,€and Leinert et al (1981) put α = 2.1.

€ Rowan-Robinson et al in this paper explored models with a modified form€Qf (β0 ) = (cos β0 ) exp(−α sin β0 ) .
A particular form of this with Q =1, α = 2.6 had been used by Murdock and Price€ (1985) to fit the ‘ZIP’ data (from Aries rockets).

€

47

1.4 Recent developments in ZL research

The cloud was assumed to be made up of isotropically scattering ‘grey’ particles (forwhich they need to be of radius greater than a millimetre), and a temperature wasassumed of the function:1

 r − 2
T (r ) = T0  
 A
No attempt was made to model for the IRAS dust bands at this point, or for the fact that€ the Sun might not lie at the centre of the symmetry plane of the disc, which had beenreported by Dermott et al (1987); but provision was made for the symmetry plane of thedust not to be coincident with the ecliptic plane, and to be inclined to it at an angle i ,with a longitude of the ascending mode Ω .
The best fit was found to be: Ω = 79º , i =1.1º , Q = 0€€
and a figure very different from Leinert’s 2.1 estimate, of α = 4.9 . This means
physically that in €this model, the dust is closely confined in the direction perpendicular
to the plane of symmetry, and not so closely confined in the direction of radial distance€from the Sun; the disc is quite flattened.

€

Rowan-Robinson et al (1991), in a following paper, used this model to subtract themain, smooth distribution of this theoretical cloud, revealing the shape of the dust
−1bands. Figure 1.8 shows the result, at 60µm . The contours are 0 − 5 MJy sr , in−1

steps of 1 MJy sr .

€
€

€

48

1.4 Recent developments in ZL research

Figure 1.8 From Rowan-Robinson et al (1991). Grey-scale and contour map of the 60 micron emission,
in Galactic coordinates, after subtraction of RHVW and application of a correction factor.

These authors concluded that there was a dust band at β = ±10º , at colour temperature
about 170ºK, and ‘warmer’ bands at β = ± 5º , barely detected at 60µm , with a colour
temperature of 250ºK. The inner band was identified by Dermott et al (1985) as being
due to debris from collisions in the Themis€
asteroid family, situated at about 3 AU. It is€ stays in the vicinity of the€ dust due to this ‘family’
proposed that the collisional
asteroids themselves, forming a diffuse toroid of material, which only slowly spirals intowards the Sun under the influence of drag from the Poynting-Robertson effect andsolar wind. Members of families of asteroids all have the same orbital eccentricity, andthe debris, at least initially, preserves this. Why are the bands in pairs? Dermott andNicholson (1989), in a numerical study, suggested that the toroidal distribution formedby collisions has a dust density “higher near the surface of the torus than within it”.
Thus, looking at a sectional view through the torus, we see a concentration of emission
at the edges. In the case of the Themis band, therefore, the torus is 5º thick from ourviewpoint, or in actual size, 0.25 AU. The outer bands were later identified as being
associated with the Eos family of asteroids, which also reside at about 3 AU, but whose
orbits are presumably more scattered in their inclinations to the ecliptic. In addition it
was noted that there could be a component of interstellar dust, forming a disc located at
50 AU, in the ‘inner’ Oort Cloud, beyond the orbit of Pluto, which is thought to be theorigin of short period comets.

49

1.4 Recent developments in ZL research

Jones and Rowan-Robinson (1993), noting that the previous model did not make goodpredictions at low ecliptic latitudes, and also did not account for the IRAS Dust Bands,
modified the ‘Fan’ empirical model at low values of z , where z is the perpendiculardistance of the dust particle from the plane of symmetry, and thus sin β0 = z . A newrmodel was constructed, using components corresponding to an inner smooth fan of€€relatively high density, an outer fan, discontinuously beginning at about 1.5 AU, and€superposed toroidal distributions corresponding to Themisand Eos. Based on the
hypothesis that ‘old’ drag-degraded asteroidal Bands are at least significant contributors
to the smooth cloud, the Fan model was assumed to contain essentially no material
beyond a distance of 3 AU from the Sun. The major source of perturbation to dust grain
orbits was assumed to be the planets. For areas close to the Sun, the parameter γ was
set to 1, since this is what the Poynting-Robertson effect indicated might be the steadystate situation, not 1.3 as used in analysing the Helios data. In this paper, the authors
now modelled the density profiles of the bands themselves, using €an analytical modelfor two families of asteroids, Eos and Themis. They also looked at modelling thedensity profiles of the broad bands seen in Figure 1.8, which the authors suggest may
be due to the ‘scattering out’ of smaller particles from the narrow bands. The model
they used is an adaptation of the narrow band model, but with two exponents instead of
one to simulate the broad ‘plateau’ in the emission profile. The authors comment that it
was this broad band which was actually first discovered by Low et al (1984). The best
fit for this broad band model was added in to the fans and the narrow bands, plus anisotropic component, to achieve an overall best fit to the thermal emission mapping. In
the fitting of the new Fan model, they found these new values for the variables:Ω = 76, i =1.3, α = 8.0, Q = 2.85 .
Figure 1.9 shows the appearance of the modified combination ‘Fan and Band’ model.

€

50

1.4 Recent developments in ZL research

Figure 1.9 From Jones and Rowan-Robinson (1993). A density map of the best-fitting combined
Fan and Band model. The grain density is normalised.

Jones and Rowan-Robinson in this analysis concluded also that the asteroid belt only
accounted for about a quarter of the dust in the main belt. The rest was presumed due tocomets, and material from the Oort Cloud, plus the possibility of an interstellar
component.
Kelsall et al (1998), modelling the COBE data, referred to Dermott et al (1996) andLeinert et al (1998), discussed the evidence from the capture of particles, and concluded
that 45% of the interplanetary dust was asteroidal. This may be compared with Durda
and Dermott (1997) who estimated that 34% was from collisional destruction ofasteroids – 10% from family members, and 24% from the main band. So it would seemfrom these analyses that perhaps one third of all interstellar dust is from asteroids. We
have seen that Baguhl et al (1996) found evidence of interstellar component, but Kelsall
estimated, from their own fitted model, that its contribution is ‘small’ (less than 1520%). The distribution model of Kelsall et al (1998) is almost exactly along the lines of
Rowan-Robinson et al (1990) except for the use of different symbols. Their formula forn(r ) in terms of a function of zc

Rc

corresponds to f (β0 ) . The parameters they

arrive at for best fit are:€€4.14, Ω = 77.7º , €i = 2.03.
γ (or Q) = 0.94, β =This paper gives a very useful picture of density contours in a slice across the plane of€symmetry (Figure 1.10), showing the appearance of dust bands and the Earth Ring.

51

1.4 Recent developments in ZL research

Figure 1.10 From Kelsall et al (1998). Isodensity contours of the IPD model components, shown for across-sectional slice perpendicular to the ecliptic plane: (a) all components combined; (b) smooth cloud;(c) dust bands; (d) circumsolar ring. Contour levels in (c) and (d) are a factor of 8 smaller.

Grogan, Dermott and Durda (2001) were the first to construct a model of the ZC dust
bands including a significant population of large particles, as suggested by theobservations of Love and Brownlee (1993). Moving away from empirical models
entirely, Grogan et al followed the journey of asteroidal debris through the Solar
System, constructing a model based physically on individual dust particle orbits, taking
into account the size-frequency distribution of the material. A given asteroid collision
produces debris according to some size-frequency distribution, given byn( D) ∝ D 2−3q ,and, for ‘collisional equilibrium’, Grogan et al put q ≈ 1.8 .

€ noted that the average inclination of the Eos family of asteroids did not exactly
Theymatch that of the ‘ten degree’ (actually€9.35 degree) dust band, but, working from thebreak-up of a single family member in the lower inclination end of the family, they
were able to produce an impressive fit, shown in Figure 1.11.

52

1.4 Recent developments in ZL research

Figure 1.11 Best-fit dust band model (dotted curves) of Grogan, Dermott and Durda (2001) to the
12 micron IRAS observations in both the leading (top) and trailing (bottom) directions of view. Thismodel has a size-frequency index of 1.43, indicating that large particles dominate the distribution.

The excess velocities of collision products will be much less than the orbital speed ofthe parent particles (15-20 km/s), so the orbits of the individual fragments will besimilar, their orbital elements only slightly perturbed in the beginning. But in only 100
to 1,000 years, this small differential in relative velocity (10-100 m/s), corresponding toa minor dispersion in semi-major axes, produces a thin ‘tube’ of material spread all
around the parent asteroid’s orbit. Over the next few thousand years, secular precession
randomises Ω , the ascending nodes of the elliptical orbits, producing a toroidal
distribution, the extremes of which are observed as dust bands.

€

53

1.4 Recent developments in ZL research

The main driver of secular precession is the gravitational effect of Jupiter, but all the5planets have an effect. Over the next 10 years, Poynting-Robertson drag lessens theeccentricity of the orbits, circularising them, and steadily reduces their semi-major axes.
Figure 1.12 shows the predictions of Grogan et al for the evolution in inclination andascending node of a typical€
debris ring as a function of time, this time from a collisionin the Koronis family. It can be seen that the end result, is a randomised I f and Ω ,such as might contribute to a main continuum, the dust simulated in “Fan” models.

€€

Figure 1.12 From Grogan, Dermott and Durda (2001). Showing the evolution of the inclination
distribution of a wave of 10 micron diameter Koronis dust particles migrating from the asteroid belt
toward the Sun (Kortenkamp and Dermott 1998). The particles were released 52,000 years ago with amean semi-major axis of 3.02 AU. The wave has now reached 1.01 AU, indicating approximately 3,500
years between each frame. The motion of the ‘centre of mass’ of the distribution represents the changing
I f and Ω of the wave.
The radius of the circular distribution remains constant until the wave approaches Earth and many
particles are gravitationally scattered. The large open circle in each frame follows the evolution of asingle orbit and€ reveals the clockwise precession of Ω p .
€The predictions of Grogan, Dermot and Durda’s models€ favour the argument that thesefeatures, the dust bands, are non-equilibrium features of the cloud, i.e. they are the result

54

1.4 Recent developments in ZL researchof individual, random, catastrophic collision events, rather than the gradual grinding
down of asteroid family members.

1.4.2.2 ISO (Infrared Space Observatory) data
In spite of its faintness in optical wavelengths, the Zodiacal Light is the dominantsource of the medium-infrared sky brightness as seen from the Earth. The cross-overpoint is at a wavelength around 3.5 microns; at wavelengths shorter than this, the majorpart of the energy is scattered sunlight (with Doppler-absorption lines) – and at longer
wavelengths the spectrum is dominated by thermal emission of dust grains. One of themost interesting parts of the whole spectrum is in the mid infrared (between 5 and 16
microns wavelength range), where the surface brightness increases by two orders ofmagnitude. In addition to the practical importance of characterising this steep rise inintensity, this is the spectral regime to look for spectral signatures of the materials
constituting the small dust grains – emission lines which can be simulated in thelaboratory, and then recognised in the observations of the ZL. The general shape of the
mid infrared spectrum is relatively well-determined now (Abraham et al 1997). Verylittle is known, however of the variation of spectral shape over the sky. Thesevariations may arise from temperature differences within the Interplanetary Dust Cloud,or from differences in dust composition and size distribution in different viewing
directions. Leinert et al (2002) analysed 29 ISO/ISOPHOT spectra distributed over theaccessible sky, in order to search for variations in the spectral shape. ISO carried two
separate instruments, ISOCAM (with a resolving power of 40, and wavelength range 516 microns), and ISOPHOT (resolving power 95, wavelength range 5.9 to 11.7
microns). At such low resolutions, Leinert et al found that the observed spectral range
was smooth and featureless, and could be fitted by black body radiation with colour
temperatures in the range of 255-300 Kelvin. The variation in temperatures wasmodelled by geometrical distribution of dust in the inner Solar System. From theirresults they concluded that the present-day interplanetary cloud is well-mixed in termsof grain composition and size distribution. Later, in 2003, Reach et al (2003), observed
the Zodiacal Light spectrum, this time with ISOCAM, in a wide range of orientations
relative to the Sun, aiming to determine whether there was a silicate feature in the ZL,
and whether there were spatial variations in its spectrum. They found radiationconsistent with a cloud of particles of sizes 10 to 100 microns, and they did detect a

55

1.4 Recent developments in ZL research‘weak feature’ which they attributed to a mixture of silicates, more dominant at low
elongations from the Sun than further out (see Figure 1.13).

Figure 1.13 From Reach et al (2003). Ratio of the observed Zodiacal Light spectra to a black bodycontinuum. The continuum was fitted using only wavelengths outside of 9-11 microns.

Reach et al were able at this point to make a comparison between our Sun’s ZC infraredemission spectrum and the emission from the first exo-zodiacal cloud discovered, Beta
Pictoris. The similarities and differences, along with comparisons also with spectrafrom comet Hale-Bopp and interstellar dust, are illustrated in their diagram, shown in
Figure 1.14. Noting the wider ‘plateau’ in the emission feature in (a), and the shape of
the peaks in (b) and (c), it may be reasonable to theorise that ‘b + c = a’, and the feature
seen here in the Zodiacal Cloud comprises elements from cometary material andinterstellar material. This supports my proposition (sections 1.4.4 and 4.11) that radial
velocity measurements may indicate a significant interstellar component in the dust
which produces the optical Zodiacal Light.

56

1.4 Recent developments in ZL research

Figure 1.14 From Reach et al (2003). Continuum-subtracted spectra of (a) the Zodiacal Light,(b) Comet Hale-Bopp (Crovisier et al 1997), (c) the Beta Pictoris disc, and (d) interstellar particles(Draine and Lee 1984).

57

1.4 Recent developments in ZL research

1.4.2.3 COBE (Cosmic Background Explorer)
In 1989, the COBE satellite was launched, following the IRAS strategy, into a Sunsynchronous Earth orbit, moving in a plane which remained almost perpendicular to thedirection of the Sun. It had two infrared detecting instruments on board: the DIRBE(Diffuse InfraRed Background Experiment), which had a 0.7º by 0.7º field of view
pointing 30º off axis, looking at 10 infrared bands in the range 1.25 to 240 microns; andthe FIRAS (Far Infrared Absolute Spectrophotometer), which had a 7º diameter field ofview, pointing along the spin axis, and was sensitive to long wave infrared up to 1,000
microns. IRAS had made the discovery of the inclination of the axis of the mainsmooth ZC to the ecliptic, plus three sets of narrow dust bands, and a circumsolar ringat 1 AU, with a concentration of dust just behind the Earth. Kelsall et al (1998), using
data from DIRBE, derived a detailed spatially varying Zodiacal spectrum in the 1.25240 micron range and mapped the complex geometrical structure of the ZC, the purpose
being to subtract the Zodiacal contribution and accurately measure the remaining
cosmic DIRB. Fixsen and Dwek (2002), combining the measurements from the two
COBE instruments, built up an annually averaged infrared spectrum of the ZC spanninga huge range of wavelengths. At the lower wavelengths they found a ZL spectrum
indicative of a black body at about 240º K; but at 150 microns, they found a sharp
break, which they attributed to a cut-off at around 30 microns radius, which for particles
predominantly composed of silicates, corresponds to a mass of 3. 3 ×10−7g. They18

estimated the total mass of the Zodiacal Cloud to be 2-11 times 10 g, requiring a
14

replenishment rate of 10

grammes per year. The spectrum of the components of€emission in the night sky which Fixsen and Dwek (2002) derive is most informative.

€

€

58

1.4 Recent developments in ZL research

Figure 1.15 From Fixsen and Dwek (2002).
Overview of the spectrum of the night sky in the entire range of the infrared.

Figure 1.15 is adapted from their figure, which shows that around the area of the breakin the Zodiacal dust spectrum, the Galactic Interstellar Medium (ISM) and the DIRB
near peak intensity, but beyond about 500 microns the CMB (cosmic microwave
background, echoes of the Big Bang) becomes completely dominant. This dominancecontinues all the way up to 5cm wavelength!

1.4.3 Collected samples and in situ measurements
Interplanetary dust particles of sizes 1 micron to 1 millimetre have been collected fromthe Stratosphere, polar ice, and from deep sea sediments (Jessberger et al 2001). By
direct laboratory examination, it has been discovered that typically the particles areblack, porous, and have a chondritic composition.

They are complex mineralaggregates, sometimes known as ‘fluffy’, a mixture of large numbers of sub-micrometer
sized components. One such particle, collected from the Stratosphere, is illustrated in
Figure 1.16.

59

1.4 Recent developments in ZL research

Figure 1.16 From Jessberger et al (2001). A porous chondritic IDP collected from the Stratosphere.

These individuals are believed to be reasonably representative of the population of
IDP’s in the vicinity of the Earth’s orbit. The smallest survive unmelted because of
their low entry speeds in the atmosphere, and only moderately altered by their capture
and residence on Earth. One feature distinguishing these finds as a group is that they
contain large abundances of noble gases, having been exposed to the solar wind andcosmic rays for a long time. Helium and Neon as well as many spallogenic isotopes arefound, confirming their origin outside the Earth. Not all of the particles are fluffy – a
few are smoother, but of similar composition. By measuring the abundances of two
different isotopes of Oxygen in such particles, Keller et al (2003), ascertained thattypically just one of the many tiny particles making up a fluffy IDP was of interstellar
origin … the so-called Stardust. Determination of the mechanism of the capture of sucha grain is a challenge! The current opinion is that the origin of IDP’s has not been
firmly ascertained.
Pioneers 8 through 1 1 , H e l i o s , Galileo, and Ulysses space probes measured
interplanetary dust in situ from 0.3 to 18 AU distance from the Sun (Grün et al 2001).
The consensus is that within 3 AU the particles orbit the Sun at low inclinations (lessthan 30 degrees), and their orbits are moderately eccentric (eccentricity less than 0.6).
The measured dust flux displays a 100-fold decrease in density from the orbit ofMercury to that of Jupiter. In addition, particles in highly eccentric orbits have been

60

1.4 Recent developments in ZL researchrecorded in the inner Solar System. The so-called Beta-Meteoroids are predicted toleave the Solar System on unbound hyperbolic orbits, accelerating because of the actionof solar radiation pressure. Where did they come from … ?

1.4.4 Interstellar Dust
One of the great ‘surprises’ in space observations, as noted by Sykes et al (2004), wasthat the detectors on board the space vehicle Ulysses, once past the orbit of Jupiter,
began to register particle impacts from the opposite direction to that expected frominterplanetary dust particles, and at high velocities, clearly indicating an interstellar
component to the dust cloud, which predominates at these distances from the Sun, but
has been now detected, by the Hiten detector, even at 1 AU (Grün et al 1993). If myinterpretation paper had been published as planned around 1975, this discovery might
have been less of a surprise! This is the source of both satisfaction and frustration tome, from the position I occupy in 2007, but a full discussion of the possibility of aninterstellar component of the ZC, mirroring my writings in 1974 based on our
measurements, will be offered in Chapter 4.
Grogan, Dermott and Gustafson (1996) estimated that the contribution of interstellar
dust flowing into the Solar System was small – but the conclusions of Landgraf et al in2000 were somewhat different. By attributing all collected particles with impact
directions not attributable to elliptically orbiting dust, and speeds in excess of the local
escape velocity to an interstellar component, they inferred a flux of interstellar grains
with masses > 10

−14

−6

−2 −1

kg of 10 m s . This is therefore an estimate of the dust

density in the Local Interstellar Cloud (LIC).

Confirmation of the interstellar

component has also come from other quarters. Particles moving in hyperbolic orbits€ in ‘AMOR’ ground-based radar measurements by Baggaley (2000)
have€been detectedof meteors, of limiting particle radius 20 microns. It also appears that these particles
have unusual origins, to say the least – Baggaley (2000) apparently saw evidence for
influx from the direction of Beta Pictoris, well-known for possessing its own ‘zodiacal
cloud’! Landgraf et al (2000), combining Radar estimates with in-situ measurements,
found a flux for interstellar particles of mass greater than 10−14

I will return to this discussion in section 4.11.1.

€€

kg , of 10 −6 m−2 s −1 .

61

1.4 Recent developments in ZL research

1.4.5 Doppler measurements and models since 1974
Doppler shift radial velocity measurement was taken up again in 1978, by Fried (1978),using a different type of instrument – a correlation mask radial velocity spectrometer.
This ingenious device, first devised and described by Griffin in 1967, correlates the
Doppler-shifted spectrum under study with an unshifted reference spectrum, using aphotographically produced optical mask, and gives a single output, recorded by a single
detector, in which the information on radial velocity content is immediately available.
Fried found no difference in the magnitudes of radial velocities found in morning andevening. He also concluded from his data that the orbital velocities of the dust particles
were in excess of those for elliptical Keplerian orbits, and that particles within 0.7 AUof the Sun were moving at speeds greater than the escape velocity (see section 4.13.1).
After my period of experimentation was completed, it was some years before our own
Doppler-Shift based line of inquiry was continued, in a 3-year experiment, by myformer supervisor Ken Reay and I. R. East, using a spectrometer of improved stability,
and their further results were published in Astronomy and Astrophysics (East and Reay
1984). They confirmed the morning-evening asymmetry found in my observations, andfound the material of the Gegenschein to be receding from us at 2.5 m/s. They found noevidence to support the results or conclusions of Fried: in particular they did not support
the idea of hyperbolic orbits. For detailed discussion, see section 4.13.2.
After another 10-year period of little activity in the radial velocity field, the subject wasrevisited by Clarke et al (1995). They presented a more rigorous theoretical model ofthe ZC, and included the full quantum mechanical scattering function of the particles, as
successfully used by Hong (1985) in purely photometric predictions, to predict thedetailed shape of the absorption line. They concluded that asymmetries should be seen
at high resolution, revealing the fine structure of radial velocities. In particular they
offered an interpretation of the F-Corona observations of Beavers et al (1980), whorecorded double minima in their absorption line spectra. No observations were
published following on from these new predictions.
Thereafter, the line of Doppler-shift experiments seems to have lapsed, until the quest
was taken up by Reynolds et al (2004). Their technique, using the Wisconsin Hydrogen

62

1.4 Recent developments in ZL research

Alpha Mapper (‘WHAM’) Fabry-Perot spectrometer attached to a 0.6 metre telescope,with a one-degree field of view, and a resolved width of 12 km/s (equivalent to 0.2 Å)
was immediately much more powerful than our small light-gathering coelostat andfilter-limited Fabry-Perot etalon. Their results have greatly reduced the error bars on
the ‘Radial Velocity (= wavelength shift) versus Elongation’ graph, and they have beenable to employ them in comparison with modern sophisticated models of the ZC withgreat effect. Interpretations of these new data were published by Ipatov et al (2006).
Using measurements with rather better signal-to-noise than we were able to achieve,and higher resolution, this team have attempted to extract other parameters from thespectral curves to supplement the radial velocity estimations; they have been the firstobservers since my experiment to measure the behaviour of the width of the line withelongation (see Chapter 4). They also obtained spectra out of the ecliptic plane; an
example is shown in Figure 1.17.

Figure 1.17 From Madsen et al (2005). Spectrum of the twilight sky (red) and the Zodiacal Light
toward the North ecliptic pole (circles), centred near the MgI 5183.6 angstrom line.

Their observations are compared with new theoretical curves calculated by Ipatov et al(2006). Their models compute actual spectra pertaining to a line of sight, and are basedon evolutionary theories of the Zodiacal Cloud (Ipatov and Mather 2005); rather than

63

1.4 Recent developments in ZL researchsimply starting with an arbitrary set of orbital elements for the dust particles, Ipatov et
al go back a step to scenario in which about 15,000 asteroidal, cometary, and transNeptunian particles begin a journey through the Solar System. They compute the pathsof these objects over time, under the influence of Poynting-Robertson drag, solar winddrag and the gravitational disruptions of the planets, and freeze the aggregate to obtain
predictions of the Doppler shifted spectrum resulting from a line of sight. One of themost important parameters which controls the dynamical behaviour of dust particles is

β (not to be confused with the ecliptic latitude parameter β0 ). β is defined as the ratioof the Sun’s radiation force to the gravitational force on the particle, and typically has
values from 0.0004 to 0.4 in these Zodiacal Cloud models.

€

It has the value

the particle diameter, and QPR (notβ = 0.573QPR /(ρs) , where ρ is the density,€s is €to be confused with Queen and Paul Rodgers) is the radiation pressure coefficient,which is close to unity for particles larger than one micron. For silicates of density

€

€
€
€
ρ = 2.5g / cm 3 the values of β 0.004, 0.05, and 0.4 correspond to particle diametersof about 120, 9.4, and 1 micron respectively. Evidently if β becomes equal to or
greater than unity, the particle will be swept out of the Solar System. This is the basis€

for the assumption€that the ZC does not contain significant numbers of particles lessthan 0.3 microns in diameter. Particles in exactly €this balanced state of force will beable to drift through the Solar System without interference from the Sun; however they
will experience perturbations from the planets! I will return to this idea in the‘continuous flow’ model discussed in Chapter 4.
I will revisit the ‘WHAM’ recent findings in Chapter 5, and will be able to compare
their findings with my measurements of 30 years earlier.
It is a remarkable circumstance for me that, after a period of relative neglect of almost
30 years, this line of inquiry, Doppler measurements of the ZL, should become active
again, at the very moment when I have been able to return to the field. In the final
chapter of this thesis I will outline current plans to take this study further in the near
future.

64

2.1 Choice of wavelength for the observations

Chapter 2
Preparations and experimental details 1971-1974
This chapter details the preparations for making the radial velocity measurements,
and the manner in which observations were actually performed. It includes the
choice of wavelength and resolving power, the site, the building of the observatory, andthe construction of the Fabry-Perot Interference Spectrometer – its optics, control
electronics, counting electronics, calibration equipment, and information output devices.
The optimisation of the instrument is then described, and finally the techniques ofoperation are detailed, in the two observing periods during which my data werecollected. In the following sections it should be noted that I inherited a system used in
the experiments of Reay and Ring, but it soon became apparent that this was subject toinherent shortcomings and unreliabilities, and in any case needed substantial changes tooperate at a new wavelength in a new, purpose-built observatory. My first year of workon the PhD was spent in redesigning and rebuilding the optics and electronics for an
essentially new Fabry-Perot system, featuring a number of important improvements,
which will become apparent below.

2.1

Choice of wavelength for the observations

Radial velocity measurements using the H β region of the spectrum had proved
difficult, owing to the presence of high rates of H β emission (Hindle et al 1968), so I
chose to study a different region: the Fraunhofer line at 5183.6 Å, present in the Sun’s
spectrum because of absorption€by relatively cool Magnesium ions in the reversal layer.
The profile of this line in pure, unshifted€ sunlight is shown in Figure 2.1. The hope wasthat this line, which is deep and relatively free from parasitic absorption by other lines
in the Sun’s spectrum, would NOT be present in emission in the spectrum of the ZL.

65

2.1 Choice of wavelength for the observations

66

2.2

2.2 Choice of resolving power

Choice of resolving power

This choice had been the result of much previous debate (Reay 1968; Smeethe 1970)
which will not be repeated here. However, the outcome may be summarised by notingthat if the resolving power of the instrument is too low, the depth of the line isunderestimated, and the exact position of the minimum point becomes inaccurate; and if
the resolving power is too great, the throughput of the instrument with such a narrowprofile becomes so small that very low light levels have to be measured. In this case,since only small numbers of individual photons are being counted, photon shot noise
becomes dominant, and measurements of the shift in position of the minimum point of
the line again become inaccurate. The optimum resolution is thus somewhere between
these extremes.
A resolved width of one and a half Å was decided upon in this case, since this gave
enough resolution to ‘see’ a well-defined absorption line, and at the same time allowed
enough light through to give adequate signal-to-noise in the counting rates. It may seem
strange to be making measurements on a line only about an angstrom wide with such awide instrumental profile, but the choice for optimum resolved width has to be madebased on the information sought, and the method used to elicit it. My data reduction
methods are detailed in Chapter 3, and it will be seen that, processed in this way, theaccuracy of the results is quite insensitive to resolving power.

2.3

The Site

The Testa Grigia observatory (in the perennial snow above Cervinia, Italy) having beenabandoned as unusable owing to unpredictable weather, an excellent site for ZL
observations was offered to us by Prof. Francisco Sanchez of the Universidad de la
Laguna, at the (then) newly founded Observatorio del Teide at Izaña, situated at analtitude of 7,700 feet on a volcanic ridge in Tenerife, in the Canary Islands (Figure 2.2).

67

2.3 The Site

Figure 2.2 Location of the observatory on the island of Tenerife.

The weather conditions of Tenerife are dominated by the Trade Winds and its huge
central dormant volcanic peak, known as El Teide. For much of the year, conditions arestable enough to produce an almost permanent inversion layer at a level well below thesite of the observatory. The resulting view from Izaña is, more often than not, abeautiful “Sea of Clouds” below, with only occasional Cirrus clouds above. ‘Seeing’ isusually very good, as first ascertained in 1853 by Charles Piazzi Smythe on his famous
expedition for site testing, and the only things which can interfere with good
observations are dust in the air (locally known as ‘Kalima’), blown in from the Sahara
desert when the prevailing winds are strong enough in a particular direction, moonlight,
and latterly some light pollution from the ever-growing English and German ‘ghetto’ of
Playa de Las Americas in the South of the island, along with some new resorts being
built along the South-East coast. In 1973, in our own experience, the Izaña skies wereclear for about 70 per cent of the time. The site is high enough, and of suitably low
latitude (28 degrees North) for the Zodiacal Light to be favourably positioned most of
the time, in a good, dark sky in both morning and evening, allowing observations to bemade at quite low elongations from the Sun. At various different times of the year, boththe morning and evening Zodiacal cones stand almost perpendicular to the horizon. The
observatory also had, and has, the advantage of being accessible directly by road from68

2.3 The Site

Santa Cruz and La Laguna, both less than an hour away. We were frequently able toobserve the ZL with the naked eye, a conspicuous and splendid sight, comparable inbrightness with the Milky Way. On nights of exceptional transparency the whole
Zodiacal Band could be traced by eye as an extension of the morning and evening
cones, a continuous faint band of light around the ecliptic, with a distinct brightening
frequently discernible at the antisolar point, rendering the Gegenschein an
unmistakeable though faint feature around midnight.

2.4

The ZL Building

Figure 2.3 The hut viewed from the South, with coelostat assembly in open position.

At the time of commencement of my experiment, the Observatorio del Teide consistedof just one Spanish ‘hut’, actually quite well appointed, built to facilitate ZL
experiments mounted by the Universidad de La Laguna. It was decided to build a new
hut to house the Imperial College observations. In collaboration with Professor Ring
and Professor Sanchez, I designed and supervised (mainly from a distance, with the helpof Carlos Sanchez on site) the building of a new ‘Coseley’ English hut to house ourexperiment (Figure 2.3). The building was modular, and specifically tailored for the ZL

69

2.4 The ZL Building

work, with a rudimentary living space, a thermally insulated laboratory, and a control
and electronics room. The building was designed to withstand winds of up to 90 miles
per hour, a necessary requirement for survival in Izaña in the Winter months. The ZL
Building was situated and oriented to give optimum sky coverage to East, South, andWest, with a two-mirror coelostat mounted at the South corner as shown in Figure 2.4.

Figure 2.4 Layout of the ZL building and coelostat.

2.5

Instrumental details

2.5.1 The Coelostat, and sky area accessible
Our interferometer was highly sensitive to movement, and it would have been
impractical to attach it to a moving telescope. So instead the device of a coelostat wasused, a kind of adjustable periscope containing two 6-inch aperture mirrors, to channel
light from any chosen area of the sky into the Fabry-Perot instrument, which remained
static inside the hut.

70

2.5 Instrumental details

Figure 2.5a The coelostat, set on concrete plinth, showing positions of its mirrors.

Figure 2.5b Showing the path of light through the instrument into the hut.

71

2.5 Instrumental details

The coelostat mount was equatorial; Declination and RA could be manually set to areas
on the celestial sphere, pre-chosen to be free from obvious contamination by brightstars, the Galactic plane or nebulosity. The field of view was approximately one degree– and the sidereal drift over the length of our scans dictated that the instrument
produced spectra from areas measuring about one degree by three degrees.
The coelostat was mounted on a concrete plinth (see Figure 2.5a) with a channel in thetop, forming a square section tube leading into the hut at a height of about four feet. A
cover was put on the outer end of the tube, hinged so as to be able to be swung aside,
providing access for lining up. In the tube was an optical bench assembly supporting athird (final) mirror to direct light from the coelostat in a horizontal beam into thelaboratory, and the final f/17 lens of the coelostat, both with kinematic mounts, tofacilitate lining up (Figure 2.5b). The whole coelostat was equipped with a metal
weather-proof housing, attached by means of a hinge, so it could be easily swung clear
for observing.
An oblique mirror in the laboratory directed the incoming beam from the coelostat on to
the entrance aperture of a gas-pressure scanned Fabry-Perot interferometer. The mirror
was rigidly attached via a kinematic mount to the interferometer itself.
The resulting sky area accessible to our instrument is shown in Figure 2.6a, and shows asuperb horizon from East to South, especially favourable to good observations of themorning ZL, which had been less well covered by previous experiments. Accessibility
to the entire visible length of the Zodiac at any time was excellent. In the diagram theeffective horizon is drawn on a projection of the sky, marked with two different coordinate grid systems: Altitude/Azimuth, and Hour Angle/Declination. This diagram isa smaller version of a chart I constructed to make quick conversions between the twocoordinate systems on site. Such a “Stereographic” projection of a spherical surface isconstructed by placing a plane surface to touch one point of the sphere, and a point
source of light at the exact opposite point on the sphere; the shadows formed on theplane surface form the projection (Figure 2.6b). This projection is convenient for
accurate geometrical work because it has the property that all circles on the sphere arerepresented by circles on the plane surface.

72

2.5 Instrumental details

Figure 2.6a Sky accessible from coelostat marked on a coordinate conversion chart.

Figure 2.6b Illustrating the stereographic projection.

2.5.2 The Fabry-Perot Interferometer (F-P)
The essence of such an instrument is that light to be analysed is passed through two
successive filters (see Figure 2.7). The Fabry-Perot etalon itself can be set up to allow anarrow band of wavelengths to pass through (one and a half angstroms, the chosen
resolved width – see Figure 2.7a). However, as shown later, used alone, the F-P etalon
would also pass a whole series of other pass-bands, separated by a wavelength interval

73

2.5 Instrumental details

known as the free spectral range. In order to suppress these unwanted wavelengths, awider band filter (about 6 Å) is used, known as a premonochromating filter. The passband of this filter is illustrated in Figure 2.7b. The result of placing the two filters inseries is shown in Figure 2.7c. As the F-P peak is scanned over the range of
wavelengths to be studied, all the other peaks are filtered out by the premonochromator.
The height of the peak transmission varies according to where it stands with respect tothe filter shape, and this effect has to be compensated for in the reduction of the data
later. The parameters shown in Figure 2.7 are roughly those chosen for this instrument.
The theory of the Fabry-Perot etalon is considered in more detail later, in section 2.6.5.

Figure 2.7 Principle of the Fabry-Perot Interferometer.

A schematic of the optical system and electronics is shown in Figure 2.8, and the gaspressure system in Figure 2.9.

Figure 2.8 Schematic of optics, and control and counting electronics.

74
2.5 Instrumental details

75

2.5 Instrumental details

Figure 2.9 The Nitrogen system.

.
The Fabry-Perot etalon was housed in a pressure pot into which Nitrogen gas was fed.
To scan the instrument, Nitrogen pressure was varied over a range of 6 atmospheres,allowing the instrument to scan through a range of 9 Å in wavelength. The equipment
incorporated a fully automatic gas pressure stepping and cycling control, the ‘pressure
controller’, with which the wavelength points to be sampled were selected on a
programmable switchboard. Detection of the transmitted light was by a speciallyselected EMI 6256 photomultiplier, whose dark current was reduced to about one persecond by cooling it to –30 degrees Celsius by means of a Peltier Effect cooler housing.
This compares with a count rate of about 5 for the night sky. Pulses from thephotomultiplier were amplified and counted by Nuclear Enterprises (Edinburgh series)
scalers and ratemeters.

76

2.5 Instrumental details

Compensation for sky transparency was provided by diverting a few per cent of theincident light through a gelatin filter with a 400 Å wide wavelength pass-band, centred
on 5200 Å, into a second channel of pulse counting electronics. In operation, thepressure controller cycle was started by a pulse from the ‘Sampler’ (described below).
The controller allowed the pressure to increase until a programmed point was reached,
and then, after a stabilising period of 5 seconds, issued a pulse to instruct the Sampler tobegin counting photons at this wavelength. The Sampler reset the scalers and allowed
them to count for a predetermined time; it then closed the gates and initiated a print-outof the total count received in each channel. When the print-out was complete, thepressure controller was automatically started off in search of the next sample point, andso on, until the end of the scan, when the controller returned the pressure pot toatmospheric pressure. Throughout the scan an analogue measure of the signal was
produced by the ratemeter, which was attached to a chart reader, so a continuous visualcheck on signal levels was available while the scan was proceeding.

2.6

New features of the instrument

In this section I detail my own improvements, in which most of the equipment wasredesigned and rebuilt.

2.6.1 Fore-optics
The front end of the F-P optics was changed so that the coelostat’s ‘final’ lens now
delivered an f/17 beam directly from the coelostat mirrors, removing the necessity for arelay lens, eliminating two glass-air surfaces and hence increasing the light throughput.

2.6.2 Field lensA field lens was added at the entrance aperture, imaging the final coelostat lens on to
the F-P plates, while the coelostat lens images an area of sky on to the entrance
aperture. This was designed to prevent anomalous results due to any non-uniformities
in the sky appearing on the plates, which themselves are not perfectly uniform. Also,
because the plates were now at a pupil of the system, maximum throughput wasachieved for a given plate area, giving increased luminosity.

77

2.6 New features of the instrument

2.6.3 Beam splitter for compensation channel
The glass plate which is used to deflect a small percentage of the light into thecompensation channel has to sample the full area of the beam, but since this channel
functions over a 400 Å wide pass-band, little intensity is needed, and in this situation it
is vital to steal as little light as possible from the signal channel. The plate was thus
coated with a single dielectric layer for maximum transmission at 5000 Å, at 45º, an
unusual specification. The thickness of the layer was calculated as below. Figure 2.11
shows the geometry.

Figure 2.10 Coating for maximum transmission.

Remembering that the glass plate is to be positioned at an angle of 45º to the incoming
beam, we have immediately, by the laws of refraction

ng sinθ g = nl sinθ l = sin 45º =12

where ng and nl are the refractive indices of the glass and the layer respectively. If the€

thickness of the layer is t , maximum transmission occurs when the path difference

€

€

€

78

2.6 New features of the instrument

between the two reflected rays 2nl t cosθ l is exactly one half wave, making theminterfere destructively:€
λ2nl t cosθ l =
2

€

€

∴t =
=

λ
1
.
4 nl cosθ lλ 1 24nl 1−  nl 2 

λ=

4nl 1−

1
2nl2

The material used was Magnesium Fluoride for which nl = 1.37, and for a λ of 5183.6€

Å, the formula gives a thickness of 0.1 µm (microns) for the layer. This, when applied
to both surfaces, reduced the percentage of diverted light to one half of that diverted by€€ the throughput of the spectrometer.
the previous plain glass plate, further increasing€
2.6.4 Construction of housing
A new light-proof cover was made for the instrument, part of which is visible in Figure
2.11, with sides which could be slid off for easy access to the optics without disturbing
the alignment of the etalon. A cam mechanism was made for tilting the broad bandinterference filter, to enable its wavelength peak to be ‘tuned’ without opening thecover. Some new lateral adjustments to the lenses made it possible to line up the wholesystem more accurately. Most importantly, the three control screws for adjusting
parallelism of the etalon were re-mounted remotely, on top of the box (Figure 2.11),
where they could be operated by an observer at the ‘Maxwellian’ view of the plates intransmission, the eye located on the optic axis by an iris diaphragm. From such a
position one can immediately see if the plates are in alignment, because the illuminationis uniform, or as close to uniform as the perfection of the polishing and coatings allows.

79

2.6 New features of the instrument

Adjusting the screws while looking at the Maxwellian view, it became easy to get a‘feel’ of the etalon, and much realignment time was saved during observations. An
improvement in the response of the adjustments was effected by replacing the metal
runners on which the plates rested by PTFE (Teflon) strips, reducing friction andlessening long period drift.

Figure 2.11 The centre part of the Fabry-Perot box, showing the etalon in its pressure pot,
and remote adjusting screws behind.

2.6.5 Fabry-Perot theory, and the finesse of a practical interferometer
The F-P plates are the heart of the spectrometer, and upon their quality depends thequality of the results. Since a new operating wavelength had been chosen, 5183.6 Å,
the opportunity was taken to have some new plates made. They were cast from fusedquartz, four inches in diameter and about an inch thick, polished to an accuracy of

λ 200 and coated with 7-layer dielectric (alternating Magnesium Fluoride and Zinc
Sulphide) stacks to give a reflectivity of 95 per cent at 5000 Å. These materials are stillused today. Some recourse to Fabry-Perot theory is necessary here to explain what€

performance was expected, and what was obtained.

80

2.6 New features of the instrument

Figure 2.12a The ‘perfect’ Fabry-Perot etalon.

Figure 2.12b Transmission function of the perfect Fabry-Perot etalon.

2.6.5.1 Reflection Finesse N R
Figure 2.12a shows a ‘perfect’ F-P - two perfectly flat plates with a high reflection
coefficient R, accurately parallel and separated by a thickness t of dielectric material€ n . Any textbook of light (e.g. Jenkins and White) shows that a rayof refractive index
€

€

81

2.6 New features of the instrumentof light incident at an angle θ to the normal emerges as an infinite number of multiplyreflected rays, and the superposition of these rays, known as the Airy Summation, (seeAppendix 7) gives a resultant transmission τ , given by€1− R)2(τ = € 2

1+ R − 2R cos Δ.........

(2.1)

where Δ is the path difference between successive transmitted rays. The geometry is
the same as €in Figure 2.10, and the path difference is given by€

Δ = 2nt cosθ .

Maxima of τ€ occur when€πΔ
= Mπ where M is an integer, called the order.
λ

π .2nt cosθ €
= Mπ
€
λ

i.e.

∴ For a maximum, 2nt cosθ = Mλmax.............

(2.2)

€
Used as an interferometer, the Fabry-Perot is worked at around zero angle of incidence,

€

so

€
Normal Incidence

2nt = Mλmax..............

(2.3)

€
A plot of τ as a functionof λ has the classic Fabry-Perot shape shown in Figure 2.12b– a system of narrow peaks in transmission, separated by regions of almost no
transmission at all. As was shown above (Figure 2.7), all but one of these peaks is€blocked by the€premonochromating filter, allowing just one narrow band of
wavelengths to pass through the system.

82

2.6 New features of the instrument

The separation of adjacent maxima, λ1 - λ2 , is called the Free Spectral Range, Δλ ,of the etalon. For these maxima, using equation 2.2,

€ €
and

€

2nt = Mλ12nt = ( M +1)λ22€2nt (λ1 − λ2 ) = λ1λ2 ≈ λso, eliminating M ,€(since λ1 and λ2 are both approximately equal to the central wavelength λ ).

€€€ ∴ Free Spectral Range

€

λ2
Δλ =
2nt

. . . . . . . .€.......

(2.4)

A constant called the reflecting finesse N R is defined as

€

€

Δλ
NR =
€
w

..................

(2.5)

where w (Figure 2.12b), is the full width at half power of the peak, also called theresolved width.€N R is a constant for a given set of plates, so this ratio is also constant.

€
€

€

NR =

π
Δλ
,
=

−1 1
w

2sin

F

Reflecting Finesse

NR =

π R€
1− R

where F =4R

(1− R)2..............

(2.6)

If the gap t is increased, or if the refractive index n of the material in the gap grows in€ gas pressure, the whole pattern of peaks in transmission moves
value due to increased
down along the wavelength scale, keeping its proportions; Δλ decreases at the same

€ rate as w because the ratio between them€ is a constant, the finesse of the plates. Hencefor a given pair of plates we can obtain better resolution only by decreasing the free

€

€

83

2.6 New features of the instrument

spectral range, which necessitates a narrower filter to cut out unwanted orders, and socuts down the range of wavelengths over which the instrument can scan. Note that, inequation 2.6, R is a quantity close to unity, so

R is even closer, and a good

approximation for π R is 3. So as a rough guide,

€

3
NR ≈
1− R

€

€
..................

(2.7)

Now N R has been carefully designated the reflecting finesse, because in a practical€

Fabry-Perot there are other factors which affect the ratio between Δλ and w .

€€2.6.5.2 Aperture finesse€

The first of these is brought about because there must be a finite angle of acceptance atthe plates in order to produce a finite intensity at the detector. In other words we are nolonger concerned with normal incidence only. For each angle of incidence a slightly
different wavelength is allowed to pass, so the result from a finite cone of light acceptedat the plates is a ‘smearing’ of the transmission peak. The greater the angle of the cone,
the wider the peak, and the more the effective finesse is degraded. An aperture finesse
is defined as the ratio of Δλ to the width δλ A of the smeared peak in the absence of
any other effect (i.e. assuming infinite reflection finesse). If i is the angle made with
the normal by the extreme ray of the cone, we have, from equation 2.2,€€
2nt cosi = Mλ1 , 1 2nt1− i 2  = Mλ1.
 2 

or if i is small …

€€

but, for the central ray,2nt = Mλ2€

nti 2∴€δλ = λ2 − λ1 =
M

€

€

84

2.6 New features of the instrument

N A = Δλ δλso the aperture finesse

Mλ2
NA =
giving, using eq. 2.4,
2n 2 t 2i 2€€
Aperture Finesse

€

NA =

A

(In which once again λ1 ≈ λ2 ≈ λ )

λ
nti 2

€
..........

(2.8)

Note that this has defined a throughput ( Φ = Aω ) of the plates where A is their area

€ of the cone.
and ω is the solid angle

€
Φ =

So

€

πD 2ω
4

where D is the diameter of the plates.

€
Throughput of F-P plates

π 2 D 2i 2
Φ =4..........

(2.9)

This, for maximum efficiency, should be matched to the coelostat’s throughput, which
can be calculated at the €entrance aperture. Here, if the diameter is d and the f/ numberof the beam is f ,πd 2 πΦ = Aω =
.
4 4f2

€

€

π 2d 2.
=16 f 2For matching, the throughputs are equal, giving …

€π 2 D 2i 2π 2d 2=
4
16 f 2€€85

2.6 New features of the instrument∴ i2 =

d2.
4D 2 f 2

Substituting this value in equation 2.8,€
NA =Aperture Finesse

4 λf 2 D 2
ntd 2..........

(2.10)

Thus the physical parameters of the optics alone determine the aperture finesse at which€ operated.
the instrument should be

2.6.5.3 Defect finesse N D€

Figure 2.13a The effect of polishing errors.

Suppose the plates are not perfectly flat but have a residual microstructure ofindentations in their surfaces. They may also have a large-scale curvature or ‘bowing’(Figure 2.13a). In either case a typical departure δ p from a plane surface can be found.
The usual method for determining the flatness is to misalign the plates slightly,
producing the appearance of parallel bands. Measurement of the deviation of a band
indicates the magnitude of a departure€from flatness. This means that some parts of the

86

2.6 New features of the instrument

plates effectively have a gap differing by δ p from that expected, again producing asmearing of the maximum transmission peak. A defect finesse is defined in the same
way as an aperture finesse, by analogy with the reflection finesse.

€

ND =

Defect finesse

ΔλwD.............

(2.11)

where wD is the ‘smeared’ width due to plate defects, in the absence of any other€effect. For the two extreme cases we have, from equation 2.3€

2nt = Mλ1()

2n n + δ p = Mλ2

andtherefore

= MwD

€
wD€=

so

2nδ p = M (λ2 − λ1)

€
2nδ p
M
€

.

Hence, using equations 2.4 and 2.11, the defect finesse

€

λ2 M
ND =
2nt.2nδ p

Defect finesse (polishing)

€

ND =

λ
2nδ p

...........

(2.12)

Thus, given a limit to the accuracy with which the plates can be polished, there is a limit€ For this reason, the defect finesse has been called by someto the finesse attainable.
writers a limiting finesse. However modern methods of polishing can give a flatness of

λ 200 or better, and it is the next consideration which tends to limit the overall
finesse.

€87

2.6 New features of the instrument

2.6.5.4 Coating non-uniformities
To obtain the very high reflection coefficients needed to give high finesses, quarterwave layers of dielectric material, of alternately high and low refractive index aredeposited in a ‘stack’ of 7, 9, or 11 layers, etc. A seven-layer stack can give an intensity
reflection coefficient of 95 per cent over a restricted spectral region (a few hundred
angstroms), whereas aluminium surfaces can reach no better than 92 per cent. The
layers are deposited in an evaporation chamber at a high temperature to give mechanicalstability to the layer, and the plate is usually rotated and moved in an epicyclic path toproduce a deposition as uniform as possible. The thickness is monitored optically, andwhen the centre reaches, say, a quarter-wave thickness, the plate is masked, awaiting thenext layer. Any residual non-uniformity is rotationally symmetrical, and, since eachprocess is similar, any error tends to be repeated in successive layers, giving acumulative effect. To this may be added the effects on the substrate (glass plate) of
heating and cooling, which, depending on rates of change of temperature and on how
well the plates were annealed initially, can give a distortion of the substrate itself. A
typical section of a non-uniformly coated plate is shown in Figure 2.13b.

Figure 2.13b The effect of coating errors.

A fourth ‘coating finesse’ could be defined, but it is more convenient to add a ‘coating
defect’ δc to the polishing defect δ p and modify the defect finesse accordingly.
Unfortunately the effect of non-uniformity in coatings may be more serious still,

€

€

88

2.6 New features of the instrument

because reflections from a stack of layers differing slightly from λ 4 in thickness
suffer a phase change which is also cumulative. This effect magnifies the purely
geometric effect by a factor of the order of 2.

€
We can say that the total defect δ is given by:δ = δ p + αδ c€.....................

(2.13)

where δ p is the polishing defect, δc is the coating defect, and α is the magnifying€

factor due to phase effects. Lissberger (1969), showed, by reducing a stack tosuccessive single-layer problems, that in a stack of k layers, small departures δ s in€€
€ optical path from one quarterwave in the s th layer, numbered from the substrate,
produce a phase change in the final reflected ray of

€

€

k+1 €

∑ Lsδss=2where Ls is the weighting factor which is greatest for layers near the gap, since the€amplitude is greatest here. For most purposes δ s is the same for each layer and so for a
k+1€ given system we can identify ∑ Ls as the magnification factor α .
s=2€

The modern method of calculation is to treat the whole€ system using field matrices, an€ itself to the use of a computer to ‘solve’ any configuration
approach which lends
directly. The method gives identical results for small errors in layers.

2.6.5.5 The choice of plates
To summarise, one goes to the Fabry-Perot shop armed with three equations 2.7, 2.10,
and 2.13 :89

2.6 New features of the instrument

3
1− R

Reflecting finesse

NR ≈Aperture finesse

4 λf 2 D 2
NA =
ntd 2

Defect finesse

ND =

€

€

λ
2n(δ p + αδ c )along with a knowledge of the resolution and scanning range required. The resulting‘real’ instrumental finesse N is a complex function of all these three, given roughly by

€

€

1
1
1
1
= 2+ 2+ 2
2N
NR NA ND

N turns out to be roughly half the value of any one component finesse when all of these
are roughly€equal. Any mismatching brings about a waste of available resolutionluminosity product. (ref. Chabbal 1958).
For our September trip, the plates, 4 inches in diameter, were polished flat to better thanδ p = λ 100 , giving a possible polishing defect finesse of 50. The 7-layer coatings
were accurate to one quarter per cent. This corresponds to a total expected error in thetwo stacks (14 layers in all) of one quarter per cent of 14 λ ; i.e. about λ

4

€

100 .

The magnification factor α for Zinc Cryolite is about 2. Hence

ND =

1

 1
1
+ 
2
€ 100 50 

≈ 17

€

€

The reflection coefficient was expected to be about 94 per cent, giving a reflection€ finesse of 50. It was thought that the estimation of defect finesse was pessimistic andthe reflection finesse optimistic, but the reverse turned out to be true. The plates
suffered some deformation in coating, degrading δ p , but the coatings gave an excellent
reflection finesse, close to that theoretically possible with 7 layers.

€

90

2.6 New features of the instrument

Figure 2.14 (top) Maxwellian view of F-P plates with effectively monochromatic illuminationin the green, showing non-uniformity, and (bottom) possible cross-section of plates (exaggerated).

Figure 2.14 shows a Maxwellian view of the plates illuminated normally with a
Mercury vapour spectral lamp green line. The band-width of the illumination was muchless than one angstrom. The photograph is a good illustration of a situation in which thereflection finesse is much better than the defect finesse, showing a microstructure ofsmall scratches swamped by a long range defect indicated by a bright ring, or annulus,
instead of the even illumination over the aperture that one might have expected. As thegas pressure was increased, the bright ring expanded, showing that the gap was wider atthe edges than in the centre; the plates were exhibiting a convex spherical surface, rather
than the perfect flat profile they were designed to have. Each small area on the plates
would give a transmission peak about one angstrom wide, but the peak arrives atdifferent pressures for different parts of the plates, so the overall peak is very wide.
Masking off the edges improves the overall finesse as shown in Figure 2.15, but at the91

2.6 New features of the instrument

expense, of course, of throughput. A two and three quarter inch mask was considered
the best compromise, and this was used in making the observations.

Figure 2.15 Finesses using a 2” entrance aperture (as distinct from aperture of F-P plates)

In this case, the excellent reflecting finesse is not only no advantage, but actually
decreases the throughput.
A full discussion of the limitations in throughput of Fabry-Perot interferometers due tomismatching of finesses was given by Jacquinot (1954). Operating our plates at lessthan full aperture significantly decreased our counting rates, and degraded the eventual
signal to noise in the scans. Thus for future similar work it may be advisable to use
fewer coatings, perhaps five, in the hope of preserving the overall finesse by matching
the contributory finesses more closely. The instrument was finally operated at a finesseof 10, a resolved width of one and a half angstroms, and a free spectral range of 15 Å,along with a filter width, to half power, of 6 Å.

92

2.6 New features of the instrument

2.6.6 Intensity calibration box

Figure 2.16 The new front end, with calibration equipment in place.
The square section coelostat tube is visible at the extreme left, where light enters the hut.

To calibrate the instrument in intensity, an attachment was made, shown in thephotograph, Figure 2.16, to the front end of the spectrometer. Its construction is shownin Figure 2.17. The source was an M/F pilot bulb, calibrated at the National PhysicalLaboratory for luminous intensity and colour temperature at a standard voltage. The
source illuminates a Magnesium Oxide screen, newly deposited, through a set of neutral
density filters. The angled screen was made big enough to over-fill the acceptance coneof the instrument, reproducing the conditions of sky operation, and the intensity wasarranged to be comparable with a typical signal level from the night sky, so that a valid
comparison could be made. The voltage was monitored on an AVO meter, and kept
constant with a rheostat. When not in use the bulb was switched off and the screen was
swung out of the beam. For details of use and calculations see Chapter 3. The intensity
calibration bulb is not visible in figure 2.16; it is housed in the calibration box, as shownin Figure 2.17. The lamp visible in Figure 2.16 is a wavelength calibration source.

93

2.6 New features of the instrument

Figure 2.17 Intensity calibration box.

2.6.7 Wavelength calibration assembly
Two spectral lamps were used to perform calibration scans. For accurate wavelength
calibration the intensity of the lamp must be constant during the time of the scan, andthe illumination of the entrance aperture must be as uniform as the signal light, sinceevidently each annulus in the entrance aperture passes at a slightly different angle
through the F-P plates, peaking at a slightly different wavelength. The former condition
is not met if the lamp has just been switched on, so this lamp was left permanently
switched on during a night’s observing. This necessitated a rig which was easily
demountable, to be hidden in a light-proof enclosure during sky scans. The assembly
located a single-layer paper diffuser at the entrance aperture, illuminated evenly by thecalibration lamp from a distance, via a 45º angled mirror.
The lamps used were a Phillips Zinc Halide, giving a spectral line at 5181.95 Å, and a
Pye hollow-cathode lamp, giving the emission line corresponding to the Fraunhofer
5183.6 Å line, which was our designated target in the spectrum of the Zodiacal Light.

94

2.6 New features of the instrument

2.6.8 Control electronics

Figure 2.18 The New Sampler module, which determines integration times
and generates the various pulses to operate the machine sequence.

A new, reliable sampler module was constructed using MDTL logic, of my own design(Figure 2.18). The unit, whose logic diagram is shown in Figure 2.19, triggers thevarious functions in the control system, and incorporates a timer which determines theduration for which counts are recorded for each sample wavelength point.
The timer was completely new, designed and constructed by myself, since the old timer
was the source of many problems. It is simple in design, but gave a new level ofaccuracy and reliability.

95

2.6 New features of the instrument

96

2.6 New features of the instrument

Instead of counting mains cycles, which can be erratic, a ‘clock’ waveform was
generated by a ring oscillator, whose frequency was governed by a single temperaturecontrolled capacitor. Rather than counting pulses and decoding at various points toachieve various time intervals, the logic was set to give a signal after a constant numberof pulses, and the time interval was varied by switching different controlling capacitors
into the clock loop, thus changing the clock pulse rate. Intervals of one fifth, five, ten,
fifty, and one hundred seconds were available. The sampler logic diagram (Figure 2.19)
uses Motorola logic symbols and most of the elements perform their standard nandfunction. The short period delays, however, are a little unusual, and warrant some
explanation.
Figure 2.20 shows a single element of a typical nand gate used as a short period delay.
(a) shows the arrangement using the logic symbol for the nand gate, and (b) is theequivalent circuit diagram, showing the individual npn transistor elements inside thedevice. The nand element is simply diode logic followed by a saturating amplifier.
When all inputs are floating at 8 volts (designated ‘up’), both transistors conduct, andthe output is 4 volts (‘down’). If only one of the inputs is held ‘down’ at 0 volts, VT1’sbase becomes 0 Volts, cutting it off; VT1’s output is thus down at 0 volts also, whichcuts off VT2, giving a delayed output of 8 volts (up).
In the delay arrangement, a negative edge is applied via a capacitor of capacitance C
directly to the base of the gate, turning the gate off and giving an ‘up’ output. The
second plate of the capacitor now charges up through the device itself, through resistors€ leakageR1 and R3 at a rate depending on ( R1 + R3 ) × C . In this case no internal
resistor is needed. When the second plate reaches 8 volts the gate is switched back,giving a delayed negative-going edge at the output. The technique is limited to a few€€ milliseconds, because use€ of a large capacitor soon over-runs (and burns out) the
previous element in the circuit; MDTL elements are essentially fast, low fan-out
devices. Figure 2.20b provides a key to the internal circuitry of the nand gates, which,prior to the invention of multiple Integrated Circuits, were the ‘state-of-the-art’elements for constructing the controllers and computers of the 1970’s.

97

2.6 New features of the instrument

Figure 2.20a Logic representation of a short period delay.

Figure 2.20b The equivalent circuit.

In operation, the sequence was initiated manually, or by a pulse from the pressure
control unit indicating that the previous sample had been completed. The scalers were
reset by a negative pulse, and after a small delay, the ‘gate’ bistable was flipped to the‘open’ position. This initiated a pulse to start the scalers counting signal events, andalso allowed clock pulses through to the four ‘sixteen’ counters in series. When 16

4clock pulses had been counted, the change in output level of the last counter was used toflip the bistable to the ‘off’ position, allowing no more pulses to be counted. A pulse€

98

2.6 New features of the instrument

was issued to close the gates in the signal scalers, and, after a short delay, pulses weredelivered to trigger the print-out sequence and instruct the pressure controller to moveto the next wavelength point. In this configuration there is fact no actual decoding of
the clock pulses needed, and the counters do not have to be reset; they are already atzero, ready for the next sample.
There is an additional advantage in this system in that an error due to faulty starting or
stopping of the clock can never give a large percentage error in the interval delineated.
4

An error of a single period gives rise to a fractional error of one in 16 . This is sosmall that no elaborate systems are needed to count an exact integral number of clock
pulses.

€
The sampler was constructed in modular form so as to be compatible with the Nuclear
Enterprises equipment (Edinburgh series).

New power supplies were madeincorporating fast voltage control cards in addition to a Claude-Lyons mechanicalvoltage stabiliser, and a Voltstat saturating transformer to minimise the disruptive effectof the highly erratic mains supply existing in Izaña in the early 1970’s. Figure 2.21
shows how the logic elements were physically arranged within the Sampler, on fourhome-made printed circuit boards, plugged in to the internal frame.

Figure 2.21 Showing layout of components on Sampler boards.

99

2.6 New features of the instrument

The pressure controller was mounted in a new easily accessible box, separate from theprogramming switches and logic, minimising vibration problems during operation.

2.6.9 Output devices
To facilitate on-the-spot scrutiny of counts recorded, a somewhat unreliable paper tapepunch was replaced by a Nuclear Enterprises print control unit and an Addo line printer.
For each spectral point the signal count and the compensation channel count wereprinted on a special roll giving duplicate copies.
The analogue outputs of the rate meters were fed into a new multi-channel chart
recorder along with the reference pulses generated by the pressure controller at each
step, providing an excellent visual display and permanent record which could bechecked for spikes and other anomalies.

2.7

Optimisation of the system

2.7.1 Optics
Much care was taken to ensure that the optics were correctly aligned. The optical benchin the channel under the coelostat was used to define a principal optic axis. The
coelostat circles were set to the Sun’s position accurately calculated from ephemeris
tables, and the coelostat frame was positioned to project the Sun’s image on to thecentre of the third mirror, which was in turn adjusted to send the image exactly along
the optic axis of the tube. The oblique mirror in the hut was then adjusted to project theimage centrally on to the entrance aperture of the Fabry-Perot etalon. The F-P optics
were already aligned by autocollimation, with the plates set perpendicular to the optic
axis by illuminating them with a bright extended source and lining up the centre of thering pattern on the exit aperture. It then only remained to make sure that the F-P axis
coincided with the axis of the incoming beam, which was centrally illuminated by theSun’s image. Using this procedure it was possible to point the coelostat to an accuracyof half a degree, more than adequate for this work.

100

2.7 Optimisation of the system

2.7.2 Counting electronics
The pulse height spectrum of a real signal differs from that arising from noise inside thephotomultiplier tube, and an optimum signal to noise figure can be obtained by correct
positioning of the threshold voltage level above which a signal pulse is counted, andcorrect setting of the EHT voltage used for the photomultiplier. The quantity to be
optimised, as related by Hindle and Ibbett (1966), isS
S + 2D
where S is the signal count and D is the ‘dark current’ count. To optimise this,€readings of counting rates with the shutter open and closed were recorded for a numberof settings of EHT and threshold levels, and the value of this formula calculated in each€ case, taking care to allow€time for the tube to settle down after any change in EHT.
2.7.3 The system as a whole
To test the equipment as a system, numerous scans of the spectral lamps were
performed, revealing the shape of the instrumental function, and how it changed invarious conditions of use. Experience was gained in this way of the system’s stability
while undergoing vibration and temperature changes, and its response to adjustments.
Drift was at its most troublesome in the early evening when the external temperature
was dropping quickly. In the first observing trip this effect was minimised by aligning
the etalon at sunset to give a Magnesium I emission line peak at a slightly higher
position in the pressure range than that eventually required, and gradually tightening theadjustment screws to bring the peak down in pressure to the working value at the timeof the first sky scan of the night. In this way the wavelength drift due to compression in
the spacers and springs was crudely balanced against the drift due to temperaturechange, and stability was easier to obtain. In the second trip using better temperature
control equipment (fan heaters with thermostats, fitted inside an enclosure made ofpolythene sheeting) the drift was largely eliminated. As well as spectral lamp scans andscans of the filter using a continuous spectrum white light source, scans were recordedof the day sky. These furnished a useful standard absorption line shape, which was

101

2.7 Optimisation of the systemcompared with a theoretical convolved shape (see Chapter 4) to check that the correct
resolution was being achieved. Test scans were also recorded of moonlight, brightstars, the Milky Way, and ‘dark’ sky, to check for reasonable counting rates to be
expected. Attention was then turned to the Zodiacal Light.

2.8

How the observations were made

Zodiacal Light measurements can only be made in the absence of twilight andmoonlight. A typical observing period begins about three days after Full Moon, when
only about an hour of complete darkness is available between the end of astronomical
twilight and moonrise. The available observation time then increases each night, up to a
maximum at astronomical New Moon, and thereafter diminishes as the crescent ‘New
Moon’ of every-day parlance sets later each evening. The observing period ends about
three days before the next Full Moon. I have represented a number of such cycles
diagrammatically in Figure 2.22, which is a useful quick guide to viable ZL observing
time available in 1971. Of course clear weather is also necessary, and an absence oflight pollution. Around New Moon both morning and evening Zodiacal Light can be
seen. Also shown on the chart are the angular elevations of the principal axes of themorning and evening cones; the October period is seen to be most suitable forobservations of the morning ZL at its maximum elevation (about 87 degrees at Izaña).
For similar reasons April is most suitable for observing the evening ZL. These periods
fortunately also give freedom from contamination by the Milky Way at low elongations
from the Sun. I planned my observations to be mainly concentrated in the ecliptic
plane, and before an observing trip a number of points along the line of the ecliptic weretabulated.
The method adopted was to choose points separated by about five degrees along theecliptic characterised by a number according to their ecliptic longitude, and note theircoordinates in Right Ascension and Declination. Points were chosen as far as possible
free of nebulosity and bright stars, but were later to be improved by experiment.

102

2.8 How the observations were made

103

2.8 How the observations were madeThe procedure for making an observation was as follows: one of the tabulated points
was chosen, its suitability verified by eye and reference to Figure 2.6, and the coelostat
was set to accept light from this point. Counting rates were then sampled for between
10 and 18 spectral elements over a range of about 4 Å centred on the 5183.6 Å line.
The integration time for each sample point was 50 seconds. On each night the points
available were systematically covered, using the periods close to twilight to look at low
elongations. From night to night, the tabulated ecliptic points gradually shifted inelongation from the Sun, as the Sun’s position changed relative to the background stars,
and so over a few days a fairly uniform sampling of elongations was attained, with adeliberate concentration on the Gegenschein.

If any tabulated point showed aconsistently erratic pulse count in the compensation channel, a better point was foundand henceforth used instead, so by experience a good set of points was built up.
The area of the sky projected on to the entrance aperture was about one square degree,
and since the duration of a scan was typically 12 minutes, the spatially resolved patch ofsky was an oval three degrees long and one degree wide.
Unless special care is taken, any drift in the Fabry-Perot spacing will result in a drift in
the position of maximum transmission during the scan, adding uncertainties to thewavelength calibration. The pressure corresponding to the position of the unshifted
5183.6 Å line was determined immediately before and after each scan, by scanning theMagnesium hollow cathode lamp. These are the sky scan’s associated finesse scans, socalled because they are also used to test whether the finesse was good enough, or not, inwhich case the plates needed realignment. Care was taken to ensure that, in thesecalibration scans, the entrance aperture was evenly illuminated and that the intensity
was constant, for reasons given above. Scans of a continuous white light source weremade before and after each night’s observations to determine the position of the filter
peak in wavelength, and at these times dark counting scans were also recorded, thesebeing performed by covering the entrance aperture rather than closing the shutters, inorder to count stray light along with the photomultiplier dark or noise current. In the
April 1972 observing session, the stabilities of both etalon and filter were greatlyoimproved by temperature-controlling the instrument to ± 0.5 Celsius. Two methods
presented themselves for improving the stability still further in subsequent experiments;

€

104

2.8 How the observations were madean optically contacted etalon with quartz spacers, or variable thickness spacers
controlled by a good servo-loop. Both systems were later realised in the Imperial
College Infrared Astronomy Department. The servo-controlled piezo Fabry-Perot,
developed by Tom Hicks, proved to be particularly rewarding in performance: in thehands of Hicks and Reay, it gave birth to a highly successful manufacturing company
devoted to the commercial production of these devices for use in research and industry,
and in 1984, one of these instruments enabled East and Reay to revisit ZL Doppler shiftmeasurements (see later, section 4.13.2), and amass data with a significantly improved
signal-to-noise ratio.
The wavelength scale was established experimentally by scanning the Magnesium andZinc lamps alternately, and in Chapter 3 the use of these scans to make this calibration
is described. Intensity calibration scans were performed using the attachment describedabove, and scanning exactly as for a sky scan. During the day some blue sky scans
were recorded for comparison.
In all, 245 scans of the Zodiacal Light and Gegenschein were obtained, plus about 20
various others, including the apex and antapex of the Sun’s motion in the local galaxy.
This completes the account of experimental details.

105

3.1 Intensity calibration

Chapter 3
Reduction of the data
This chapter contains an account of the processes applied to the raw scans, after
the observations were completed, to transform them into a set of corrected spectra,
suitable for the extraction of parameters to compare with theory, together withsome preliminary observations on the data. I describe here the absolute calibrationof intensity, and the manner in which the scale and datum position for wavelengths were
determined, and how the scans were corrected for various imperfections, interferences
and systematic errors.

The unexpected discovery of Magnesium I emission is

discussed, along with its effect on the interpretation of the data for radial velocity
determinations. A system of averaging the fully corrected absorption line scans,grouped according to position in the ecliptic, is described, together with various
attempts to fit curves to the experimental points. The method of extracting parametersfrom the finally chosen Gaussian fits is detailed.

3.1

Intensity calibration

3.1.1 Method and calculations
The source used for intensity calibration is described in Chapter 2. The energy input tothe interferometer is calculated as follows.
The National Physical Laboratory supplies tables (Jones 1970) of Q(λ ) , the radiant
spectral intensity of a tungsten source at colour temperature Tc °K inside a glass
envelope. The value I have used in the calculations is a compromise between this figure€

and that appropriate to a black body, since for a coiled element the inside parts behave
like a black body.

€

106

3.1 Intensity calibration

Figure 3.1 Photometry of intensity calibration.

The energy per unit wavelength interval radiated by the lamp is given in the tables as:

Q(λ ) microwatts per steradian per nanometre.
Hence the spectral irradiance on the MgO (Magnesium Oxide) screen, placed at an€ angle of 45 degrees, at a distance of D cm, modified by the transmission factor τ
the filters is (referring to Figure 3.1)

Q(λ ) ×

€

1
2D

€

€
−2
−1
× τ F µ watts cm nm
2

€

€

€

F of

107

3.1 Intensity calibration

The entrance aperture accepts an f/15 beam, so a small aperture would ‘see’ on the
diffusing plate, an elliptical shape, of areaπ l 
 
4 15 

2

2 cm

2

where l is the distance of the screen from the entrance aperture. However the area

€actually seen is€much larger than this because the entrance aperture has a finite radius r

cm, so the projected shape has an extra rim of width r cm (Figure 3.1). The total area
seen is

€

€
2
π l
2
 + 2r  × 2 cm

4 15

€

Emission from such a screen (assumed non-absorbing) at an angle θ to the normal has€€an intensity per solid angleEω cosθ 2π , where E is the total incident intensity, andω is the solid angle, in this case, that subtended by the entrance aperture at the screen,
which is 4 r€

F-P isl 2 . Thus the total flux received by the€
€
€

2
π l
4r 2
×
τ
×
+
2r
2
×


F

4 15
2D 2
2πl 2 2Q( λ )

€

€

2

=

2
Q(λ )τ F r 2  l−1 € + 2r  µ watts nm .
2 2
4 πD l 2 15µ watts nm−1€

For a given finesse, a given photon counting rate represents a definite flux per unit€
€ wavelength range. A scan of the €calibration source gives a measured N

cal counts per

second at the filter. If the filter shape is normalised to unity maximum for the purposeof dividing it out from the scans, for any filter–corrected count per second N measured€
from the sky, the corresponding energy per unit wavelength
range is (now moreconveniently expressed per Angstrom):108

3.1 Intensity calibration

2 τ F
Q( λ ) r 2  lNW = + 2r  ×
 10
N cal 4 πD 2l 2 2 15µ watts/Å …..

(3.1)

This represents the flux from one square degree€ of sky received by a six-inch diameter€ aperture at the bottom of the atmosphere.

In practice there was some noise on the calibration scan, so it was divided by the filtertransmission, point by point, using a noiseless filter profile, and an average corrected
count was found for the nominal 50-second integration time.
Putting in the instrumental parameters, experimental values, and those supplied by the
NPL, for the narrow region around 5183.6 Å, the value of W arrived at was

W = 1.78 ×10 −9 µ watts per Å per count.

€
This may be conveniently expressed in photons. The energy E of a photon of€ frequency ν is given by€
E = hν , where h , Planck’s constant, is approximately 6.6261×10 −34 m 2 kg /sec .

€From the wave equation,

€

€

€

c = νλ8where c is the speed of light, approximately 2.9979 ×10 m/sec, and λ is in this case€5183.6 Å, or 5183.6 ×10 −10 m.
€€

Thus the number of photons represented by W is€€
1.78 ×10 −9 ×10 −6 × 5183.6 ×10 −10W=€E
6.6261×10 −34 × 2.9979 ×10 8

€

€

109

Hence

3.1 Intensity calibration

1 count represents 4640 photons/Å

……….

(3.2)

It is sometimes useful to know how many photons are arriving in total for a given countrate. The equivalent width of the F-P pass-band was typically 1.5 Å, so
1 count represents 6960 photons

……….

(3.3)

3.1.2 Accuracy of intensity calibration
Figure 3.2 shows a graph of counting rates measured in the laboratory using the lamp at
different voltages. An uncertainty in the AVO meter reading of 0.1 volts produces anuncertainty in counting rate of only 2 per cent. The National Physical Laboratory’s
calibration was to an accuracy of 2 per cent also, so the overall accuracy is limited onlyby photon shot-noise, small temperature fluctuations of the lamp, geometricaluncertainties and the estimation of the transmission of the coelostat. For sources
outside the atmosphere, extinction losses also contribute to the uncertainty. The final
calibration is believed to be accurate to a factor of 4.

Figure 3.2 Experimental graph of counting rate against voltage across bulb for theintensity calibration set-up.

110

3.2 Wavelength calibration

3.2 Wavelength calibration
The aim of wavelength calibration is to assign a true wavelength to every programmed
pressure point found by the pressure controller. Since in our case each programmed
point corresponded to the positioning of a rotating arm illuminating a particular pealamp (see Hindle and Reay 1967), the programmed point scale is referred to by ‘lamp
number’, L . The calibration involves (a) determination of the unshifted MgI 5183.62 Å
position for each scan and (b) determination of the separation of this line and the Zinc
5181.99 Å line in terms of ‘lamp number’. From this can be calculated the wavelength€interval between each programmed point (in ‘angstroms per lamp’). We will see that
this quantity dλdL may vary slightly with lamp number itself, because of the non-linearity of the pressure controller.

€
3.2.1 Determination of the datum line position
For the purpose of determining the 5183 Å position, celestial sources (planets, Moon)
were considered unsuitable because of their lack of availability near the time of a scan,and the difficulty of evenly illuminating the entrance aperture from a small source on
the sky. Instead a Magnesium lamp was used. Because of the slight drift of the etalon
over the period of a scan, the lamp was scanned immediately before and after each scan,and the datum line position was determined as the mean value of the two measurements.
Some early calibrations were done using a stack of diffusers illuminated quite closely
by the lamp. At the time, the arrangement was considered adequate, but soon it wassuspected that the illumination of the diffuser was slightly greater at the centre than at
the edges. To test if this was so, scans of this source were compared with scans of thesame lamp using the more evenly illuminated diffuser described in Chapter 2. The test
involved measurement of a differential of only about half a ‘lamp’ in a situation where
the drift of the peak from scan to scan and even during one scan was of the same order.
A ‘dynamic’ method was adopted, in which the drift and the differential could be easily
identified, and measured separately. The technique was to align the etalon, and thenmake a number of scans alternately using the two configurations, without realignment.
A typical result is shown (Figure 3.3), plotted as a function of time. Lines joining

111

3.2 Wavelength calibration

successive determinations of any one configuration clearly show the drift of the etalon.
The separation of the two lines at any time gives an accurate measure of the error in theold configuration, which was found indeed to be 0.5 lamps, to the blue, as would be
expected. This correction was applied to the early calibrations, and all subsequentcalibrations were performed using the new uniform diffuser for illumination.

Figure 3.3 The drift method applied to determination of error due to non-uniform
illumination of entrance aperture.

3.2.2 Establishment of wavelength scale
The separation in lamp number of the MgI and Zn lines ΔL was found by a similar
method – i.e. scanning the two lamps alternately as long as the finesse stayed good; thepositions of the peaks were plotted against time, and the separation at an instantmeasured from the graph (Figure 3.4).

€

112

3.2 Wavelength calibration

Figure 3.4 Determination of the separation of the two calibration emission lines
using the drift method.

In this case the two lines appear at different wavelength positions under the filter profile
and it would seem at first sight that the filter should be divided out to give the true
positions. However the shape of the maxima is entirely instrumental in origin, andarises from a line of virtually zero width, so the shape of the filter profile could have noeffect on the position of the peak. Strictly speaking it is only correct to divide out thefilter after deconvolving the instrumental function from the spectrum. In the case of a
monochromatic source this reduces the spectrum to a line. However the sky spectra
under study contain no important narrow features, so the filter division process is valid
for these.
With all corrections made, the values of ΔL were plotted against the position in lamp
number of the Zn line (Figure 3.5), which varied depending on how the etalon wasaligned at the time.

€

113

3.2 Wavelength calibration

Figure 3.5 Experimental graph of separation of Zinc and Magnesium lamps,
as position of Zn peak is varied.

A linear relationship between lamp number and pressure would give simply the straightline ΔL = constant, but experimentally this was not found to be the case, indicating a
non-linearity in the pressure controller. A check was made on readings of the Bourdonpressure gauge for each programmed point and a non-linearity was confirmed (Figure€3.6). For one section of its movement, the controller gave less pressure variation perlamp than expected – around LZn = 20.

€

Figure 3.6 Experimental variation of pressure as measured by Bourdon gauge
against pressure controller position in lamp number.

114

3.2 Wavelength calibration

The part least affected was around LZn ≈ 40.5, so this was chosen for future sky scans,and is a convenient point to evaluate ΔL in terms of wavelength.

€
For LZn ≈ 40.5, the €measuredvalue of ΔL was 12.9 lamps. Referring to the€pressure/lamps graph (Figure 3.6), at L = 40.5, a ΔL of 12.9 corresponds to a pressure
difference of 10.0 lbs per square inch. It is this pressure difference which is a constant€€ € for a given wavelength interval, so
going back to Figure 3.6, we can predict what ΔL
will be for any position of€
the Zn line. €The curve so obtained is shown dotted in Figure
3.5 a good fit to the experimental data. Thus a value of dλdL can€be confidently

adopted for each position on the scale of lamps. The values adopted were:€€

For L < 30dλ€dL = .134 Å per lamp........

(3.4)

For 30 ≤ L ≤ 41dλdL = .102 Å per lamp........

(3.5)

≤ 95
For 41 ≤ L€

dλdL = .136 Å per lamp........

(3.6)€

Hence a wavelength scale was established for all scans. The accuracy of the scale is the€€ the dotted curve in Figure 3.5 can be fitted to the experimental points
accuracy to which– i.e. about 0.6 lamps, corresponding to between 0.06 and 0.8 Å. The overall accuracyof the position of features in the final corrected spectra is dealt with in section 4.4.

3.3 Treatment of a single scanA typical ‘raw data’ scan is shown in Figure 3.7a, which is plotted directly from theprintout values of the signal channel. In the diagram the scale of lamp number has
already been converted to a wavelength scale, as was done for each scan using theresults of section 3.2. The arrow indicates the position of 5183.6 Å as determined fromthe two associated finesse scans of the Magnesium lamp, before and after.

115

3.3 Treatment of a single scanFigure 3.7 Illustrating the treatment of the data.

Figure 3.7b shows the corresponding compensation channel record, and Figure 3.7cplots the filter shape, an average of scans of a tungsten source performed at the
beginning and end of the night’s observing. Starting with data (a), the dark count( ≈ 60) was subtracted from each point, and the resulting true signal count divided by (b)and by (c), giving the required compensation in overall light level. The result is shown

€

116

3.3 Treatment of a single scanin Figure 3.7d. Figure 3.7e shows a typical averaged spectrum, the result of combining
the data from several scans of the same region of sky (discussed in Chapter 4).
The spectra were multiplied by a constant in each case, which was adapted to produce a‘normalisation’ taking into account the various changes in ‘comp’ channel geometry
and fore-optics during the experiments. This enabled a direct comparison to be madebetween the intensities of scans taken at different times, sometimes showing up unusual
atmospheric conditions. The fall-off at the edges of the raw data scan is seen to havebeen entirely due to the interference filter, and the Poissonian error bars areconsequently large for these ‘edge’ points.
The whole process of plotting the two filter scans for each night with their associatedfinesse scans, lining up the calibration positions and averaging the two scans, and thendetermining from this the correct filter transmission to be applied to each point in every
scan, was laboriously done by hand for the October data, but for the April results theprocess was streamlined using the experience already gained of the instrument’s
behaviour, and a computer program written by T. R. Hicks was used for the reduction
process.
Each scan was checked against comments written at the time of the scan for any
anomalies which could be assigned to local interference, or obvious failure of thecounting system. In spite of all precautions (Chapter 2), occasionally spikes in theelectrical supply lines, or aerial pickup from the nearby radio transmitter, caused thecounters to spontaneously reset. The scans were scrutinised for such effects, and in theevent of a count being lost, the chart recorder record was used to estimate a true countfor the single point affected. Only a small degradation of signal accuracy was producedby this process. The scans were inspected for any obvious asymmetry in continuum
levels. Where this could be traced to a probable misplacement of the filter in dividing
out, a better estimate was made, and the scan re-computed. Sometimes it was
concluded that asymmetries were due to other causes (pre-saturation of the signal orcomp channel during a calibration scan, contamination by clouds, stars, Milky Way or
twilight, abnormal finesse drifts or degenerations). Separate corrections for these weremade where possible.

117

3.4 Daylight correction

3.4 Daylight correction
In the case of scans recorded near dawn or dusk, a new method was adopted tocompensate for daylight contamination. In morning observations, the onset of twilight(dawn) within 10 degrees of the horizon, where much of the work at low elongations
was done, did not always coincide with the beginning of astronomical twilight(calculated as the time when a seventh magnitude star should be just visible in the
Zenith); it varied as much as half an hour depending on the local atmospheric conditions
towards the East. Similar comments apply to the end of twilight at the beginning of thenight. A scan recorded at dawn is shown in Figure 3.8. Part (a) shows the signal counts
recorded, and in part (b), the compensation channel record, the beginning of daylight
contamination, as opposed to the slight slope cause by the rising Zodiacal Light cone, isclearly visible, and it is reasonable to suppose that whole of the shaded portion isdaylight. Scans of sunlight were obtained during the day, and so this spectrum is wellknown and virtually noiseless, owing to the high daytime fluxes, making it possible toreliably subtract it from the raw spectrum, to give a much improved final compensatedscan.
If, at any spectral point, the fraction of daylight present is F (see Figure 3.8), theappropriate correction to be subtracted to eliminate daylight is F × D(λ ) × C , whereD(λ ) is the daylight spectrum scaled to unity continuum, and C is the assumed true€
continuum level in the spectrum under study. The function F × D(λ ) × C is shown in€€

Figure 3.8c and the corrected spectrum in Figure 3.8d. Apart from restoring the balance€of the continuum levels, the correction ‘fills in’ a broad area near zero shift position,€ here. The small magnitude of thesince a depression in continuum is being ‘subtracted’corrections in the important central region indicates that the method could be used togood effect even further into twilight, and at even smaller elongations. A small amountof extra noise is introduced owing to the fact that irregularities in the compensation
channel may be due to variations either in ZL flux, or the daylight flux, or both. Thus
the error bars towards the high wavelength side of Figure 3.8d are larger than normal.

118

3.4 Daylight correction

Figure 3.8 Correction for daylight contamination around dawn.
The fraction of daylight present at any point in the spectrum is p/q in (b).

119

3.4 Daylight correction

At the stage when each individual scan had been fully reduced, a preliminary
examination of them immediately revealed the existence of what appeared to be small
emission component at the MgI wavelength, variable in intensity.
This emission is treated here first, as part of the reduction section, since anunderstanding of its relevance is necessary before proceeding to the determination ofradial velocities in the dust cloud; the rest of this chapter is devoted to the MgI‘emission core’.

3.5

MgI emission

3.5.1 Discovery
In describing the central feature, the word ‘core’ is used because, whereas the main
absorption line is roughly the width of the solar Fraunhofer line, slightly widened by the
instrumental profile - about 3 Å in all – the emissions generally have widths of the orderof the instrumental resolved width itself – about 1.5 Å. They appear as a narrow peakinside a trough. An example is shown in Figure 3.9a. Figure 3.9b shows a scan inwhich there is apparently no emission.
In general, such individual scans were not of good enough quality to makemeasurements from, but for some of the data procured in adjacent parts of the sky on
the same night it was possible to construct average scans which yielded good quality
data (Figure 3.9c), and in some of these the core appeared to be real. The method of
averaging was applied in a different way for investigating the absorption feature itself,
and is described in full later (section 3.6).

120

3.5 MgI emission

Figure 3.9 Evidence for emission in the data.

It was at first thought that the core could be the result of seeing two absorption lines at a
narrow separation, giving a double feature. This can easily be shown not to be apossible explanation, however. Starting with the solar Fraunhofer line and working

121

3.5 MgI emission

from models of co-and contra-rotating dust in solar orbits, it is possible to produce a
double feature with approximately the required spacing, but the central core turns out to
be too far from the zero-shift position to fit our data, which always show the core atzero position. Furthermore, when the theoretical spectrum is convolved with an
instrumental function one and a half angstroms wide, as exhibited by our instrument,
the double feature completely disappears, replaced by a single absorption feature with aslightly greater width. It is therefore with some confidence that we believe the core to
be due to emission in the Earth’s atmosphere.

3.5.2 Wavelength of the MgI emission
In the individual scans the wavelength of the emission peak could be estimated to better
than 0.2 Å from those scans which clearly showed a core.

By combining themeasurements from all such scans the emission wavelength was determined to be
5183.6 Å ± 0.1 Å. Reference to the MIT wavelength tables shows that lines of Cobalt,
Europium, Zirconium and Magnesium occur very near to 5183.6 Å. We identify the33emission as the 5183.6 Å line resulting from the 3p P2 − 4s S transition of MgI.

€
3.5.3 Intensity of the emission€
On each scan an estimate of the height of the emission peak was made by interpolating
the probable uncontaminated absorption feature, and measuring the central discrepancy,
in normalised scale units (see section 3.3). To convert this measure into intensity it isfirst necessary to ‘un-correct’ for the compensation channel division (an average value
for the scan was used for this), and for the normalising constant, to get back to theoriginal counts, corrected only for filter transmission.
To obtain the absolute value of the number of photons arriving at the bottom of theatmosphere, the intensity calibration (equations 3.2 and 3.3) was applied, and the results
subsequently converted to Rayleighs. As an example, in the October results, the overall
average of the average nightly counts recorded in the sample period of 48 seconds
estimated to be due to emission was 17 counts.

122

3.5 MgI emission

Now since one count represents 6960 photons (equation 3.3), the number of photons
received in sample period was 118,320. The area of the six-inch entrance aperture isπ2(6 × 2.54) sq cm = 175 sq cm
4
Therefore number of photons per square centimetre in 1 second is approximately€
118320≈ 14
175 × 48
The instrument accepts a solid€ angle of 0.00056 steradians. Thus the number ofRayleighs is

€

14.08

4π

5.6 ×10

10 6

×
−4

≈ 0.32 Rayleighs.

0.32 Rayleighs is thus a typical MgI emission rate as observed. Measurements in this€€ region give extinctions of typically 0.15 magnitudes in the zenith for the
wavelength
Izaña site.
Figure 3.10 shows the average counting rates for each night for the periods of theobservations.

123

3.5 MgI emission

Figure 3.10 Mean night-time emission as a function of time. The vertical arrows indicate upper limits,and the gap in data corresponds to the period of Full Moon.

3.5.4 Variation of daily average emission
The well-established enhancement of CaII emission at times of high meteor activity(Vallance and Jones 1956) suggested to us that there might be a correlation between
MgI emission and meteor showers. The Orionid shower has a maximum activity
around October 20th and the Lyrid shower about April 21st. The September and October
1971 data show little significant variation in emission from night to night, having amean level of 0.31 Rayleighs. The April 1972 data show a generally lower emission

124

3.5 MgI emission

rate of about 0.1 Rayleighs, with some interesting structure beginning on the night ofApril 14/15, coinciding with the onset of the Lyrid meteor shower.
Two scans recorded near morning twilight on April 15th showed an unusually strong
emission of about 0.7 Rayleighs. This declined over the next few nights until anotherenhancement occurred on April 20/21, corresponding to the time of maximum activityof the Lyrid shower. Unfortunately our data do not extend beyond April 21st. It is
worth noting that the ‘enhancements’ in late April 1972 merely bring the emission rate
up to the September-October 1971 level. In this respect it is not clear whether theemission really underwent an enhancement in late April or is merely in a depressed statein early April! We believe the enhancement to be real, but the difference in ambient
emission rates September-October 1971 and April 1972 may be an artefact of thecalibration technique, and modifications to the instrument’s fore-optics between
observation periods, though every effort was made to correct for these factors.
Enhancement of the emission at the time of the Lyrid shower suggests that at least partof the atmospheric magnesium is of meteoric origin. No such enhancement occurred
during the October 1971 Orionid shower, however.

3.5.5 Variation of emission throughout the night
Depending on which mechanism produces the Magnesium emission, different variationsof intensity throughout a night could result, so an attempt was made to determine fromthe data an hourly estimate of intensity averaged over a number of nights. The process
was attempted only for the September-October period, since a greater number of scans
were affected.
The result is shown in Figure 3.11. The error bar has been calculated from the probableerrors in the individual estimates assuming a normal distribution of these about a meanvalue. Each point represents the average of about 20 measurements from individual
scans. It is seen that there is a significant increase in intensity near twilight in morningand evening. There is present in the observations, however, a strong observational
selection effect. Most of the scans recorded near sunset or sunrise were at low altitude,

125

3.5 MgI emission

and airglow phenomena are known to increase in brightness near the horizon (the Van
Rijn effect) so these scans might be expected to show a greater intensity of emission
than those around midnight.

Figure 3.11 Hourly averages of MgI emission intensity throughout the night.

Figure 3.12 Simple model of an airglow layer.

A simple model of an airglow effect is shown in Figure 3.12. If the emitting layer is
optically thin, the total intensity of emission recorded on the ground is proportional to
the line-of-sight length of path in the emitting layer, and thus …

126

3.5 MgI emission

Emission ∝ sec θ , where θ is the zenith distance.
The length of absorbing path is also roughly proportional to sec θ , but more precisely€

is proportional€to f (θ ) , a function defined using the air mass figures tabulated inAllen’s Astrophysical Quantities (paragraph 54). The intensity transmitted is thus€€
Asec θ × e−kf where k is a constant for a given air type, and A a constant governedby the physical mechanism of the emission, and determinable by experiment. If kf issmall …

€€

kf 2
intensity = Asec θ (1− kf +
)
2

€

€

The approximation ( kf small) is good only as long as the minimum transmission is€ is just realised at the lowest elongations under discussion. Using theabout 80%, whichtabulated values, the relative strengths of emission at two different altitudes can befound.

€

To predict what the effect of observational selection might have been, the scansproviding the information for points A and B in Figure 3.11 were individually checked
for zenith distance. For each scan the relative strengths from a typical airglow
distribution were calculated as above, and the average values were adjusted accordingly.
The result was that, to compensate for the observational selection effect due to zenith
distance, the value at point B needed to be multiplied by 1.44. As can be seen, thisbrings point B to roughly the same level as point A, and suggests that the true variationof intensity throughout a night is small.

3.5.6 Final comments on MgI emission
The fact that the MgI emission does not disappear late in the night when the 90-100 km
region of the Earth’s atmosphere, where most other metallic ions are observed (Narcisi
1968; Narcisi and Bailey 1965) is completely in shadow, makes it unlikely that any
fluorescence mechanism is a source, and indeed it is hard to see how fluorescence could

127

3.5 MgI emission

33

excite the triplet 3p P2 − 4s S transition. Radiative recombination of Mg+ with anelectron, the mechanism suggested by Anderson and Bath (1971) governing the daytime
balance between Mg and Mg+, would seem to be the most likely cause of the emission.

€ be expected that other Mg lines, such as the other lines of the triplet, 5172.68
It might
Å, and 5167.33 Å, would be present also. The interecombination line at 4571.1 Åconnecting the 3p to the 1s state should also be present. I know of no observational
searches for these lines, though more recent studies seem to show that neutral Mg
emission is highly variable over a day, with a maximum at dusk (Gardner et al 1995).
The discovery of this MgI emission was the subject of the first Hicks, May, and Reay
paper, published in Nature in 1972, reproduced in Appendix 5.
Attention was now turned back to the absorption line, at 5183.62 Å, for the purposes ofdetermining Doppler shifts as originally planned, and for this purpose the small
emission core, once subtracted, was ignored.

3.6

Averaging

The process of using the observed absorption line spectra to discriminate between
theoretical dust cloud models is complex. Even if we confine our attentions to a ZL
produced exclusively by a circumsolar cloud of material, with no other components,
there are still many variables. The spectral profile of the absorption line may be
expected to depend upon elongation from the Sun, and possibly the position of the Sunon the celestial sphere (reflecting the Earth’s position with respect to the Zodiacal
Cloud). There may also be short, medium and long period variations due to solar
activity, or due to changes in the density and composition of the cloud with time,perhaps brought about by the visits of comets to the inner Solar System. Because of theway in which the data was collected, it was considered practical to approach theanalysis purely as a function of elongation, but to keep the two sets of observations(from 1971 and 1972) separate, in order to highlight any differences between the two
periods.

128

3.6 Averaging

The ecliptic plane was divided into 10-degree intervals (or 5 degrees for low
elongations), and an average spectrum was computed for each group of scans in theinterval. The method was to set up a number of wavelength ‘channels’ of width one‘lamp’, and for each channel, add up the relevant number of counts recorded for eachscan, finally dividing the total by the number of scans. Care was taken to avoid
introducing artificial noise by averaging over a different number of scans for different
points. Where no value existed in a given channel, an interpolated value was used.
Because of the finite width of a channel, this method produced a final spectrum which
has been convolved with a ‘box-car’ function, one lamp in width. However the effect of
this is quite negligible compared with that of the instrumental function width.
A typical averaged scan, in this case an average of 8 individual scans, is shown in
Figure 3.7e. Most of these derived spectra exhibit what we might call a ‘well-behaved’
absorption feature. All averages were done on the central four-and-a-half angstroms of
the scans only. Attention was then turned towards fitting curves to the experimental
points, in order to extract parameters which can be compared with theory. I attempted
first to fit simple polynomial shapes to the points.

3.7

Polynomial fits

3.7.1 General
In 1973, the sole computing power available at Imperial College for processing data
was contained in a single ‘main frame’ Atlas computer, housed in the MechanicalEngineering building, and taking up a whole floor of space. Input was by means ofpunched cards or 5-hole paper tape, and due to the demand, and the slow processing
speed, it was normally possible to run a programme just once or twice a day. It was
thus a very time-consuming business to ‘debug’ a new program, and much patience wasneeded. It was possible to output on a printer, but the choice could be made to also
output on card or paper tape media, and, in this form, the result could be used directly as
input to a subsequent program. This was the approach I used, producing decks of cardsof the averaged scans and fitted curves for input later to a graph-plotting programme.
Beginning with no assumptions about the shape of the line, I computed a best-fit

129

3.7 Polynomial fits

polynomial curve for each averaged scan. The computer program used, LSQFIT, wasan adaptation of a CERNUP library program which fits, by the method of orthogonal
polynomials, a polynomial of any degree to a set of points. The fit minimises the sumof the squares of the residuals, and may be used with suitable weighting according to
the standard error on each experimental point. In practice, weighting was found to
make very little difference to the fits, and in the end, equal weighting was used in allcases.
Depending on input format, the program printed out various information on the fits.
One of the most useful is the “F-Ratio”, which gives a measure of how good a fit apolynomial of a given order is, relative to the that of a polynomial of the next order up.
As would be expected from the shape of the curves, in most cases a fourth orderpolynomial gave a good fit (three possible turning points). Figure 3.13 shows a scan forwhich it was particularly hard to find simple fit; it provides a good illustration of theway the shape of the curve can vary with successive orders of the polynomial.
Figure 3.13, as others in this report, is from a Calcomp plot. Program LSQFIT was
equipped with a card output giving the ‘x’ and ‘y’ values of the experimental points,
and coefficients of the fitted polynomial. A program POLYPT was written to accept
these as input, and plot the experimental points on the correct axes, and to trace out thefitted curve, calculated at one-twentieth angstrom intervals from the coefficients.
(Printouts of the programs used are shown in the Appendices at the end of this thesis).
These curves, fitted to the data, give an immediate eyeball estimate of the shiftexhibited by the absorption line profile, as compared with the profile in direct sunlight.
But, to remove human error from these estimates, methods were devised to extract anaverage shift parameter in each case. In this context, the word ‘average’ means a best
estimate of the shift of the line profile. From Zodiacal Cloud models, as we shall see, acorresponding ‘average shift’ can be computed, the average of the shifts of the various‘cells’ of dust particles along the line of sight integral for a particular elongation, forcomparison with this observational ‘average shift’ (see Figure 1.3). I used two separate
methods for extracting a best estimate of the shift from the fitted polynomial curves:

130

3.7 Polynomial fits

1) A determination of the minimum point of the curves (program MINIPT).
2) A determination of the wavelength half-way between the wavelengths of thetwo inflexion points (program FLEXPT).

Figure 3.13 A ‘difficult’ scan. Least-squares polynomials of order 1 to 8, fitted by LSQFIT to theaveraged spectrum for 135º West (morning), September-October 1971.

It would also be possible to make an estimate of the shift by finding an ordinate whichbisected the area contained between the spectrum cure and an estimated continuum
level.

131

3.7 Polynomial fits

3.7.2 Computation of the minimum point (Program MINIPT)
Consider a fourth order polynomial. If the curve is represented by

I (λ ) = φ 0 + φ1λ + φ 2 λ2 + φ 3λ3 + φ 4 λ4where λ is reckoned from a zero position at 5183.6 Å, and the φ ’s are the coefficients€ determined by LSQFIT, then€€∂I
= φ1 + 2φ 2 λ + 3φ 3λ2 + 4φ 4 λ3
∂λ€which must be zero at the minimum point. Since all the I (λ ) curves have the generalshape of an absorption line slightly shifted from zero position, we can assume that theturning point we want is the one closest to λ = 0 . So any value close to zero can be€ (Newton’s method).
used as a first approximation in an iterative process€ first approximation ‘root’ to the true root, we canReferring to Figure 3.14, given any
find a better approximation by finding where the tangent to the curve at this pointcrosses the zero line.

∂I∂ 2Itanθ = 2 = ∂λ ,∂λ delta
delta being the correction to be subtracted from root to give the value of ‘new root’.

€
Therefore∂I
new root = root −

€

2

∂ I

∂λ
∂λ2

132

3.7 Polynomial fits

and the whole process is repeated, using ‘new root’ as the new starting point, and so on,
iteratively. We can stop the iterative process when delta ≤ 0.001 Å.
Now, differentiating again,€∂ 2I= 2φ 2 + 6φ 3λ +12φ 4 λ2
2
∂λSo, using these values, the minimum point can be found for each curve, the distance of€ these from the zero position giving the shift directly. See Appendix 3 for the program
MINIPT used to implement this operation. The figure below is included mainly as areminder to myself!

Figure 3.14 Illustrating Newton’s method of successive approximation to solve
the equation ∂I ∂λ = 0 , giving the minimum point of I (λ ) .

€

€

133

3.7 Polynomial fits

3.7.3 Computation of the shift by the ‘inflexion points’ method(Program FLEXPT)
Again the analytical form of the 4th degree polynomial gives a simple expression for theshift measured by this method.
We require for an inflexion point

∂ 2I
=0
∂λ2
i.e.

2φ 2 + 6φ 3λ +12φ 4 λ2 = 0€This is a quadratic equation possessing roots λ1 and λ2 , and the midpoint of the line

€ joining the inflexion points has a λ value of half the sum of the roots.

i.e.

€λ + λ€1
2=€ Shift =2€=−(6φ 3 )2(12φ 4 )

φ3
.
4φ€
4

Programs MINIPT and POLYPT were designed to read directly the output cards of
LSQFIT, and the values€of wavelength shift obtained were plotted against elongationfrom the Sun. The shifts computed from the minima in the 4th order polynomial fits areshown in Figure 3.15, separately for the two observing periods.
Since 4th order and 6th order fits were available for all averaged scans, the minimum
point method was also applied to the 6th order fits, with less success because of thegreater risk of ambiguity in the turning points.

134

3.7 Polynomial fits

Figure 3.15 Wavelength shift estimates from polynomial fits. Plots of line shift vs. elongation using theminimum points of the 4th degree polynomials fitted to the averaged data as a measure of the shifted
position. This method is considered unsatisfactory (see text), but it is pleasing to note that the general
disposition of the points is quite similar to that found by the Gaussian fits method (see Figure 4.5).

135

3.7 Polynomial fits

3.7.4 Assessment of these methods
A necessary constraint on any method used to estimate the wavelength of the absorption
feature is that it use the information from ALL the points in the optimum way, bearing
in mind their various probable errors. It is also essential that the measure be identifiablewith a physical quantity calculable from models of the dust cloud. In practice thisquantity is the ‘mean shift’〈Δλ 〉 =∫ Δλ.dB∫ dB.......

(3.7)

where Δλ is the shift due to an element of brightness dB , and the integral is along a€ sight (Bandermann and Wolstencroft 1969; Reay 1969). If the original solarline ofspectrum is symmetrical (a good approximation at 1.5 Å resolution) it is easy to show

€ that Δλ is also in this case identical to the shift€ of the bisector of the inflexion points,
and to the minimum point.

€

Noise in a practical spectrum reduces the reliability of all three methods so far
mentioned. The minimum point method depends critically on the accuracy of a few
central points, so just a small emission feature, as was sometimes present, can beenough to upset the measure. The inflexion method is sensitive to the outer, more
noisy, parts of the spectrum. The method of areas would rely on a good estimate of thecontinuum level, which again depends strongly on the outer points, which are noisy. A
small error in the placing of the filter can produce an apparent asymmetry in thecontinuum levels to each side of the absorption feature. The results of the methods inpractice bear out these comments. The scatter on them all is large, obscuring even thegeneral trends visible to the eye in a casual inspection of the scans, and some results are
obviously misleading if reference is made to the curves. For the final analysis, all ofthese methods were therefore rejected, except for the special cases (see section 4.1)
where Gaussian fits could not be effected.

136

3.8

3.8 Gaussian curve fitting

Gaussian curve fitting (Program GAUSSN)

3.8.1 Computation
The method finally adopted was to fit a Gaussian curve[

I (λ ) = I1 − I 2 exp −k (λ − λ m )2]

where I1 is identifiable with the continuum level, I 2 is a measure of the line depth, k€ is related to the width of the line, and λ€m is the wavelength of the central minimum

point of the line. These four parameters correspond to the variables V1, V2 , V3 , and€€and in Hicks, May and Reay (1974). Once the Gaussian
V4 in program GAUSSN, €curve is fitted to the observations, λm is the most useful, since the difference between€ € €this and the unshifted wavelength, 5183.6 Angstroms, will give us the effective mean€

wavelength shift 〈Δλ 〉 due to the Doppler effect.
€
An adaptation of a general least-squares fit subroutine was used to make program€ compute the fits. The program used an iterative method in the fourGAUSSNdimensions, V1 ,V2 ,V3 ,V4 , and worked from a first approximation, put in byinspection. It accepted directly the same card input format as LSQFIT. The output
cards produced by GAUSSN were used by program GAUSPT to plot the points, and thefitted€ Gaussian curves.

3.8.2 Experimental case for the use of Gaussian fits
This method of Gaussian fitting is justified experimentally in that the fits obtained are
not highly sensitive to the presence of a small emission core, nor to edge effects; also inmost cases the essential residuals are small. Every experimental point contributes to thefit, and the resulting estimation of wavelength shift.
Further experimental endorsement of the method comes from the fact that the day scans
performed with this instrument exhibit a shape which fits very well to a Gaussian curve.

137

3.8 Gaussian curve fitting

3.8.3 Theoretical justification for the use of Gaussian fits
In order to simulate from theory the spectrum of the Zodiacal Light, I have used spectra
computed from dust cloud models. The aim here is to show that by mathematicallysimulating the modifications made to a spectrum by the interferometer and subsequent
reduction methods, a curve is arrived at which is very similar in form to a Gaussian.
Curves derived from a variety of dust cloud models are discussed in detail in Chapter 4,
but the method used in all cases is the line-of-sight integration through the dust cloud,described earlier and illustrated in Figure 1.3, in which it is assumed that there is aparticle number density n with the dependencen(a,r ) = ka− p r −2€
where a is the particle radius, r is the particle’s distance from the Sun, and k is a

€ constant.
€ assumed to be€ For this test, the elongation€
was chosen to be 30 degrees and the dust wasmoving in circular orbits in a prograde sense. A value of p =1 was used.
The program CONVL (Appendix 4) read in the spectrum at one twenty-fifth angstrom€ intervals, and assumed a continuum
intervals, interpolated to one fiftieth angstrom
outside the absorption feature. It then read in the filter shape (approximated by a muchwider Gaussian function) into an array, and shifted the array so that the peak wasexactly at 5183.6 Å (equivalent to tilting the filter to ‘tune’ it in the experiment). The
input spectrum was multiplied by this shape. A Fabry-Perot profile was then generated
with the same finesse and inter-order spacing as the etalon used in our experiment. The
profile was shifted to give a peak at zero wavelength shift, and then the filter-modified
spectrum was convolved over a thousand points with the F-P profile, simulating thescanning operation. Finally the resulting spectrum was divided by a filter profile as in
the reduction (widened approximately by the instrumental resolved width). The result isshown in Figure 3.16. To this curve was fitted a Gaussian in the usual way, also shownin the figure, and it can be seen that the correspondence is remarkably good.

138

3.8 Gaussian curve fitting

Figure 3.16 Comparison of a Gaussian curve with theoretical profiles.
(a) A 30º elongation profile derived from a prograde dust cloud model with circular orbits.
The dust is assumed to fall in number density away from the Sun inversely with radius of orbit.
(b) This profile convolved with the instrumental profile of our Fabry-Perot.
(c) A comparison Gaussian curve.

It is not difficult to show that for a fairly wide range of initial absorption line shapes,
and even for close double features such as might be expected if the dust is moving inboth prograde and retrograde orbits, the result as seen by the instrument is still very
nearly Gaussian in shape. Thus the use of fitted Gaussian curves to measure theparameters of the spectra not only provides an unbiased mathematical assessment morereliable than estimates by eye, but represents a very close approximation to the physicalreality. My program, GAUSSN, for computing a least-squares fit Gaussian curve to the

139

3.8 Gaussian curve fitting

experimental points, is shown in Appendix 2. GAUSSN was used in each case to
output the parameters arrived at in the fitting process - new and more accurate estimatesof the parameters of the absorption feature for all the averaged scans. The complete setof these averaged scans is shown at the beginning of Chapter 4, plotted, each with aGaussian curve fitted, by program GAUSPT, except in two cases where no good fit
could be computed; in these cases polynomial curves have been used, as noted.
At this stage, the observations were fully reduced, and in a form suitable for comparison
with theoretical models of the Zodiacal Dust Cloud.

140

4.1 Gallery of Spectra

Chapter 4Interpretation of results in terms of physical models
In this chapter, I present the fully compensated averaged scans obtained in the two
observing periods, 1971 and 1972, for all observed elongations, each with a fitted
Gaussian curve, except as noted. From parameters derived from these fits, plots ofaverage wavelength shift, line width and line depth against elongation are
constructed. These observations, the Hicks, May and Reay (HMR) data, arecompared with previous and all subsequent observations since 1974, and examinedwith reference to the predictions of various models. The models are based on
various distributions of dust moving in prograde and retrograde solar orbits, and a newmodel is presented of interstellar material flowing through the Solar System. A
geocentric component, and the effects of distant emitting matter are also considered.
New theoretical models by various authors since 1974 are examined and compared withthe data. Theories are advanced as to why there might be seasonal variations in theradial velocities measured in the Zodiacal Cloud (ZC).

4.1

Gallery of Spectra

The following 10 pages contain reproductions of the whole set of averaged spectra,
fully compensated for filter shape, overall light levels, and daylight contamination.
They show intensity plotted against wavelength for all targeted elongations, in the two
observing periods, September-October 1971, and April 1972. For ease of comparison,
the two sets of data are matched side by side according to elongation. For most spectra,a Gaussian curve was satisfactorily fitted by the method of least squares, as described in
Chapter 3, and is shown in the relevant graph. In four cases, where the continuum
levels were very different on each side, no Gaussian fit could be obtained, and for these,a polynomial fit was used instead.

Some scans were performed over a wider

wavelength range than 5 Å, but for consistency only the central 5 Å were used for thisanalysis. In any case, the extreme outer wavelengths were observed at low transmissionof the broad interference filter, and, hence, being subject to large random errors, would
not have improved the fits.

141

4.1 Gallery of Spectra

142

4.1 Gallery of Spectra

143

4.1 Gallery of Spectra

144

4.1 Gallery of Spectra

145

4.1 Gallery of Spectra

146

4.1 Gallery of Spectra

147

4.1 Gallery of Spectra

148

4.1 Gallery of Spectra

149

4.1 Gallery of Spectra

150

4.1 Gallery of Spectra

Figure 4.0 Example of an unrealistically wide fitted Gaussian curve,in this case corresponding to ‘rogue’ point (a) in Fig 4.1.

151

4.2

4.2 General comments on the data

General comments on the data

All the averaged experimental scans show a major absorption feature in theneighbourhood of 5183.6 Å. In three cases, there is a hint of a second small absorptionline, towards the red. A curious property of the data is that the scans tend to be of lessoverall height at large elongations. These scans have been compensated for variation in
the overall level of light intensity, as monitored by the compensation channel, so thisinequality is not the result of variations in intensity around the MgI wavelength. It
seems likely that the ‘comp’ channel, with its relatively wide pass-band green filter, is
picking up airglow lines, light pollution and faint stars, which, when the ZL intensity islow, significantly inflate the compensation channel counts. This reduces the apparent
height of the corrected spectrum. At the lower elongations these contributions are
swamped by ZL continuum, and the corrected height rises to its ‘normal’ value. An
additional slight rise in intensity at the position of the Gegenschein is also explainable
by this reasoning. In future this problem could be avoided by using a narrow filter forthe comp channel, closer to the width of the pre-monochromating filter of the F-P.
The results of plotting the wavelength shifts from these fits, against elongation, areshown in Figure 4.1. There is quite a large scatter in the shift values in both observing
periods, greater than that suggested by the error bars. In the case of September-October
1971, this may be the result of an unstable Fabry-Perot. However, in April 1972, thefinesse stability had been significantly improved, and we might have expected asmoother curve. In fact, the points for April do follow a curve quite well, within theerror bar predictions, except for a few ‘rogue’ points (a, b, c in Figure 4.0). Reference
to the widths of the Gaussian fits in Figure 4.2 shows a very clear grouping of thewidths into a cluster around the median (about 2Å), and a few well outside this width. I
think it is quite possible that this unrealistic width in a few of the fits indicates faulty
curve fitting. Figure 4.0 shows such a scan. If a median-width Gaussian fit is made to
the central points only (the fringe points may be even more noisy than realised at thetime, as noted in section 4.4, owing to cosmic rays), in this case at least, a much smaller
predicted shift seems likely. This suggests that it might be worth re-reducing these datawith a prior bound set on the Gaussian width, and that this would reduce the number ofoutliers. For the purposes of this thesis, the ‘raw’ output shifts of program GAUSSN
are used.

152

4.3

4.3 Information obtainable from the Gaussian fits

Information obtainable from the Gaussian fits

Figure 4.1 The Hicks, May and Reay data (HMR) - average wavelength shifts, determined fromGaussian fits, plotted as a function of elongation from the Sun. ‘Rogue’ points a, b, and c may be theresult of failure in curve-fitting (see section 4.2).

153

4.3 Information obtainable from the Gaussian fits

Figure 4.1, then, shows the result of plotting the mean wavelength shift 〈Δλ 〉 , against
elongation, for the October and April periods. In view of the possibility of Zodiacal
Cloud geometry and density being seasonably variable, the two sets of data are kept

€ (1968) and James andseparate. On each diagram, the previous results of Reay and Ring
Smeethe (1970) are included for comparison. The broken curve drawn through thepoints in each case is a third order polynomial fitted to these new data. The polynomial
fits were made using a slightly modified version of LSQFIT and were weighted
inversely as the square of the standard error as calculated for each point by program
GAUSSN. This graph (Figure 4.1), showing wavelength shift as a function ofelongation, supplemented with various theoretical curves with which my data iscompared, is the basis for most of my subsequent interpretation in this study. It isalso the basis on which all other studies of radial velocities in the ZL to date may best
be compared.
Before we look at this graph in detail, two other experimental measurements, the line
depth and line width of the absorption feature, can be extracted from parameters of thefitted Gaussian curve, for comparison with theory. The full width of the line at its ‘halfpower’ points is related to V3 as follows:[I (λ ) = I1 − I 2 exp −k (λ − λ m )2€]

is at a minimum when λ = λm€

I min = I1 − I 2

So …

€

We require the values of λ which bring the function half-way down to the value€

I
I1 − 2 2€
so for these points

€

154

4.3 Information obtainable from the Gaussian fitse€−k ( λ −λm )

2

= 12

∴ − k (λ − λ m )2 = − log e 2∴ λ − λm =€log e 2k

Now λ − λm is just the half-width required, so€

Full width to half power = 2

€

€

=

log e 2k

2.77260k

Again, referring to the Gaussian function, an exact measure of the percentage depth of€the line (residual intensity) is evidentlyI% depth = 2 ×100
I1
In Figure 4.2 are plotted the widths of the lines calculated from k as described above,again as a function €of elongation; Figure 4.3 shows the percentage depths in the same
way.

€

155

4.3 Information obtainable from the Gaussian fits

Figure 4.2 Experimental line widths computed from Gaussian fits
– plotted as a function of elongation.

Figure 4.3 Experimental line depths of the absorption feature as apercentage of continuum level, plotted against elongation.

156

4.4

4.4 Accuracy of the fits

Accuracy of the fits

The program GAUSSN computes confidence limits for all four parameters, which
depend on the scatter of the experimental points and on how closely the spectra
approximate to a Gaussian shape. However the errors on the points are not all purely
due to random shot noise. A serious source of error arises in the outside points. These
are already subject to the largest shot noise, due to their small counts, but if the filter
shape is only slightly misplaced in wavelength, in the dividing out process, these points
are seriously affected, giving an overall asymmetry to the spectrum. To estimate theeffect of this, the filter was deliberately misplaced on some test scans by a wavelengthinterval corresponding to the maximum drift of the etalon during a scan. To these false
results the Gaussian curves were fitted as before, and the variation in parameter λm ,and therefore of the ‘average shift’, was found to be about 0.2 Å maximum. This, then,is the estimate of the error bar, shown in Figure 4.1. Similar considerations apply to theerror bars in Figure 4.2 and Figure 4.3. As noted in section 4.2, there is€more scatter inthe points than the error bars would lead one to expect. In the 1971 observations, thismay be due to an F-P which was not temperature stabilised, and whose finesse was notreliable. In the April observations, with an improved F-P, I now believe that cosmic
rays, not easily distinguishable from ZL signal, may have made the peripheral points inthe scans even more noisy, sporadically, than previously believed, leading to some false
Gaussian fits, sometimes signalled in the graphs by unrealistic widths.
It is worth noting that asymmetry in itself cannot be assumed to indicate a misplacementof the filter. The effect can also be the result of not sampling over a wide enough range
to reach the continuum level, coupled perhaps with a chance low placing of the end
points. It could also be the result of a blend in the absorption feature at the ‘low’ end,
only partially resolved. Subsequent to my researches, Clarke et al (1996) predicted thepossibility of real pronounced asymmetries resulting from pure prograde models withelliptical orbits and a more complete scattering theory; effectively two shifts, due to two‘cells’ at opposite ends of the line of sight, are superposed. Also James (1969) was able
to predict asymmetries by including retrograde as well as prograde particles. Clarke’s
predictions would be noticeable only at lower elongations than were observed in thisstudy, and we will return to the question of whether the asymmetries in the HMR are
significant or not, in discussing the findings of Madsen et al (2006), in section 4.13.3.

157

4.5

4.5 Evidence from line widths and depths

Evidence from line widths and depths

The predictions of Reay (1969) were used for a preliminary comparison with my data.
Reay’s programs were used to simulate spectra from a typical prograde circular modelwith α =1, p =1, and a model similar but consisting of 50 per cent prograde material
and 50 per cent retrograde. At the resolution used in the experiment, the two lines one
might expect from the second model were not resolved; the result was a curve similar in€€ to that due to a pure prograde model (i.e. roughly Gaussian) but slightly wider
shapeand less deep. The predicted figures, for elongation 90 degrees, areWidth (Å)

Depth (%)

1.86

30.4

Prograde

2.12

26.5

Prograde and Retrograde

The difference is about 10 per cent in each case. Our experimental values for theseparameters are shown in Figures 4.2 and 4.3. The scatter on the points is quite large, as
would be expected, and a difference of 10 per cent would not show up well in these
conditions. The HMR results, as can be seen in Figure 4.2, may favour the prograde
situation, but it is evident that more and better measurements are needed to make areliable choice between models. Overall, the widths are similar to the theoreticalestimates, but the relative depths tend to be slightly greater than the estimates. This isprobably a reflection of the fact, already suspected from asymmetries, that the range ofscanning did not satisfactorily reach the true continuum level, producing an underestimate of its intensity, and hence a higher percentage depth. The possibility of either
diffuse ‘wings’ to the line, or of blends, cannot be excluded, though at this point thenarrowness of the pre-monochromating filter makes firm conclusions difficult. The
only other observers to publish measurements of line-widths are Madsen et al (2006).
Their results are shown for comparison in Figure 4.4. The resolution in theseobservations is much better (12km/s, compared with about 90km/sec), and they also
conclude that no dependence of line-width on elongation can be established, but it isvery interesting to note that the general pattern of points plotted is quite similar to ours
(Figure 4.3). The full-width at half maximum (FWHM) expected for dust in typicalasteroidal orbits (about plus and minus 20 km/sec, giving a maximum dispersion of 40
km/sec) is significantly less than the 77 km/s of their, and our, observations.

158

4.5 Evidence from line widths and depths

Figure 4.4 From Madsen et al (2006). Line widths as a function of elongation. The coloured lines aremodels which trace the motion of populations of dust particles of various origins, subject to gravity,
radiation pressure, and drag forces, with a ratio of radiation pressure to gravitational force β = 0.02.

Madsen et al (2006) remark that to fit the observations, we need
€ to consider a large
fraction of non-asteroidal, possibly cometary particles. Future work at even higher
resolutions may reveal exactly what this pattern is. Madsen et al (2006) comment that aflat bottom to their profiles, particularly in the Gegenschein area, where one might
expect to see a pure daylight spectrum, indicates large numbers of particles in elliptic
orbits. However the theoretical curve for elliptical orbits in Figure 4.4 does not fit theirobservations, or ours. The main conclusion of Ipatov et al (2006) is that asteroids alone
cannot account for the WHAM observations of radial velocities, and that particles
produced by comets, including high-eccentricity comets such as Comet 2P/Encke andlong-period comets, are needed. The mean eccentricity of Zodiacal dust particles thatbest fit the WHAM observations is estimated to be about 0.5.

159

4.5 Evidence from line widths and depths

Evidently any inferences from these data on the extent of ‘filling in’ of the Fraunhoferline by the superposition of a very blurred spectrum (Ingham 1963) would also be
unreliable, and in general, though these figures are included for the sake ofcompleteness, I have not drawn any firm conclusions from them.

4.6

Analysis of the wavelength shift versus elongation data

Since the large probable errors in the estimates of line-depth and line-width make it
difficult to analyse these graphs with confidence, we will turn our attentions now to theplots of wavelength shift as a function of elongation (Figure 4.1), for which the analysis
is much more promising. The dotted curves in this figure show a preliminary leastsquares fitted polynomial, computed for each set of data, in a program adapted from the
LSQFIT. Whilst the April 1972 data curve is fairly symmetric about the anti-solar
point, a distinct evening-morning asymmetry appears in the September-October 1971odata. The evening radial velocity maximum at ε =110 is significantly higher than theomorning maximum at ε = 40 .

€Here follow various attempts to synthesize the experimental data using dynamical
models of €the supposed Solar System dust cloud and its environs. From each model,
the behaviour of the average shift 〈Δλ 〉 with elongation is predicted and comparedwith the shifts extracted from my data. For a more thorough comparison, a full
spectrum over the range of interest could be predicted for each elongation, sampled at a

€ of the data. This kind of analysis is deferred to a future
rate comparable with thatexperiment.

4.7

Forces experienced by particles in the Zodiacal CloudIn the models discussed here, the forces acting on a particle are considered to be gravity,
radiation pressure, and the Poynting-Robertson effect (Robertson 1937). Forces due toelectromagnetic drag, the action of the Sun’s magnetic field on a charged particle, areneglected in this first analysis, which is of a quasi-steady-state velocity pattern.

160

4.7 Forces experienced by particles in the Zodiacal CloudThis decision needs some justification. Parker (1964) calculated that U/V and soft Xrays might produce a charge of +10 Volts on a particle in orbit in the Solar System. The
Sun’s magnetic field, that of a dipole, but with the lines of force massively stretched out
along the ecliptic plane by the Solar Wind, reverses in sign every few days for a particle
in orbit, and the subsequent Lorentz force on sub-micron particles tends merely todisperse them slowly from the invariable plane.
The constant drag force has been estimated as a small force comparable with the
Poynting-Robertson force by Singer (1967), and in a co-rotating plasma may be a small
accelerating force.
Belton (1966), in a fuller analysis of all the charging and discharging processes,
concludes that the force due to the equilibrium charge per particle at 1 AU from the Sun
is of the same order of magnitude as both radiation pressure and gravity for a particle of
0.1 microns radius, although he admits to a lack of information especially with regard todischarge, which depends strongly on the shape of the particle. However, Belton has
underestimated the radiation pressure at 0.1 microns by using the geometrical crosssection, so the effect of the Lorentz force is over-estimated; in any case the latterdecreases sharply with particle radius, so for a particle of 0.3 microns it is
proportionally nearly an order of magnitude less, even on this reckoning.
Thus the typical scattering time of these forces is long, and the forces do notsignificantly modify the momentary trajectory of a particle, even though they play avital part in the evolution of the interplanetary medium.

4.8

Rotating dust cloud models

Many programs have been written to predict the Doppler effects of radial velocities on
the absorption line spectrum generated by various rotating dust cloud models. Some
recent simulations will be discussed in this chapter, but for an initial assessment, thepredictions of Reay (1969) were used. Predictions were made from prograde andretrograde models of dust in orbit around the Sun. The orbits are assumed to be
circular, and the Poynting-Robertson effect is for this first analysis ignored.

161

4.8 Rotating dust cloud models

The model assumes a spatial density which falls off as r−α

at a distance r from theSun, and a distribution of particle radii a given by:

n(a)da = ka− p da .

€€€ on a particle by sunlight is proportional to the flux at the
The radiation pressure exerted€ of the particle. The force it experiences is therefore, like the gravitational
locationforce, inversely proportional to the square of its distance from the Sun, but in the
opposite direction. The effect of radiation pressure is thus to reduce the effective
gravitational mass of the Sun for the dust particle, reducing its orbital velocity for agiven distance from the Sun. But the force is also proportional to the surface area
offered, i.e. to the square of its radius, whereas the gravitational force is proportional toits cube, so radiation pressure becomes more important the smaller the particle is.
Hence we can expect the magnitudes of the wavelength shift in the predicted curves tovary strongly with p , and on whether the orbits are prograde or retrograde. In Reay’s
treatment, both reflected and diffracted components of the scattered light are evaluated,
and an integral along the line of sight gives an effective wavelength shift€〈Δλ 〉 =∫ Δλ.dB .
∫ dB

For the geometry of this method, which follows that of Ingham, see section 1.2.5 and€ 1.5; and a fuller description of the line-of-sight integral process, as applied to a
Figurecontinuous flow model, is given in section 4.11.
Figure 4.5 shows the results of Reay’s calculations for α = 2 . In both periods ofobservations the agreement at low elongations (less than 70 degrees) between the newexperimental points and a prograde model with a high value of p look promising. For€ of the right general shape at all highthe April results, the agreement continues to beelongations, but for the September-October results, the model clearly predicts a curve at
odds with the Hicks, May and Reay (HMR) observations,€ which indicate a cross-over at
about 80 degrees elongation on the morning side, as opposed to the symmetrical crossover at 180 degrees predicted in the curves.

162

4.8 Rotating dust cloud models

Figure 4.5 The Hicks, May and Reay data of wavelength shift vs. elongation, with Reay’s predictedcurves for prograde and retrograde circumsolar dust clouds, moving in elliptical orbits. Solid lines
represent models with particles of size 0.2 to 100 microns; broken curves are theresult of eliminating submicron particles, and dotted curves are for models
in which diffraction and radiation pressure are ignored.

It is noticeable that the value p = 5 in the prograde predictions seems to be rather too
small for both sets of results. This indicates that we are dealing primarily with the verysmallest particles, around 0.3 microns, assuming a density of 3 grams per cubic
centimetre, which €can exist in orbit around the Sun. As this size is approached fromabove, the repulsive force due to radiation pressure increases to the point where theeffective solar attracting mass approaches zero. A particle which is thus ‘just’ in orbit
has a low orbital speed. In other words for particles of such a size, prograde and163

4.8 Rotating dust cloud models

retrograde models become more and more similar to models in which the dust ismoving slowly towards or away from the Sun, or drifting en masse through the Solar
System. A linear drift model is described separately in 4.11.

4.9

Other considerations in rotating models

This discussion of the rotating dust cloud model has so far focussed on adjusting the pparameter to fit the results. Other parameters may also produce a change in shape of thecurves. The number density was assumed to be proportional to r−α

, where r is the€

distance from the Sun, and α given the value 2, for the purpose of constructing Figure4.5. This is a reasonable value based on measurements of intensity and polarisation as a

€ and Wolstencroftfunction of elongation, made by many investigators. €Bandermann(1969) show the€results of calculations for other values of α (the parameter is called βin this paper). The effect of decreasing α is to raise the curve, giving higher values of€〈Δλ 〉 for elongations in the range 75 o ≤ ε ≤ 180 o , so a model could be envisaged€€
which did not have a predominance of sub-micron particles ( p = 0 or 1, say) and a€number density staying constant with increasing distance from the Sun ( α = 0 ). This€latter is almost certainly physically untrue, however. Although some authors have

€ the asteroid belt, all modernsuggested that the density may increase again towards€ density decreasing
observations, including ISO (Reach et al 2003), point to a number
with radius from the Sun, as far as, and some way beyond, the Earth. The difference
between the curves for α =1, 2, or 3 is not very great, and increasing α beyond thisopoint makes almost no difference. The region close to the Sun, ε ≤ 70 , is notsensitive to radial distribution and it is in this region that the€rotating model fits most€satisfactorily, so the conclusion that small particles, a micron or less in radius, dominate
the cloud is still valid.

€

Bandermann and Wolstencroft also give predictions for elliptical rather than circularorbits. Even for eccentricities near to 1, their curves are not changed as much as forvariations in α , and such high eccentricities are highly physically unlikely. And
elliptical orbits in any case degenerate to circular under the influence of the Poynting-

€

164

4.9 Other considerations in rotating models

Robertson effect. Thus the possibility of elliptical orbits does not materially affect thevalidity of the conclusions already drawn here. However, we will see that Ipatov et al
in 2006 came to different conclusions, based on a model in which the dust is being
continuously replaced, and there will be much more to say about rotating models in
sections 4.13 and 4.14. Meanwhile let us look at other possible models, which include
no dust in orbit around the Sun at all.

4.10 Geocentric Dust Cloud (GDC) model
A thorough theoretical investigation by Shapiro, Lautman, and Colombo (1967)
asserted the absence of a mechanism by which the Earth could acquire a dust cloud of
the proportions necessary to make an appreciable contribution to the Zodiacal Light. In
spite of this, and the discrediting by Nilsson (1966) of the satellite measurements whichgave support to the speculation, the GDC theory still had vocal advocates in 1970.

Vanysek and Harwit (1970) suggested that Reay and Ring’s results for ε ≤ 70orepresent the dominance of dust near the Earth. The predictions for a model based onterrestrial dust only are shown in Figure 4.6. The HMR results certainly seem to rule€out the hypothesis, as can be seen from the high shifts in the observationsin the region
75 o ≤ ε ≤ 120 o , both in morning and evening.

€

165

4.10 Geocentric Dust Cloud (GDC) model

Figure 4.6 Predictions of Vanysek and Harwit compared with both sets of HMR data. The shaded areas
indicate the allowed range of wavelength shifts assuming a geocentric dust cloud.

It is worth commenting that Divari’s model (1965) of a cloud of terrestrial dust particles
in highly elliptical prograde orbits with the major axes in line with the Sun could
produce a blue shift in the morning and red in the evening; but it certainly could notaccount for the reversal in sense of the shift at low elongations which the consensus ofradial velocity measurements clearly shows to exist. It is my opinion, based on theabove, that cis-terrestrial dust is not a major contributor to the Zodiacal Light.

166

4.11 A uniform dust flow model

4.11 A uniform dust flow model

Figure 4.7 Illustrating the movement of the Sun and the Earth through the local interstellar medium.

It was my feeling in 1974 that a component of the ZC might well be interstellar dust,
drifting through the Solar System, due to the Sun’s proper motion through the local
interstellar medium. I believe that this was the first time such a theory was proposed,
previous discussions having been centred on asteroidal and cometal origins of the dust.
It was my intention to publish this view in a further paper in MNRAS, but unfortunately
my work at Imperial College ceased before this could be achieved. I was recently,
however, pleased to find vindication in the Ulysses and Galileo satellite dust collection
experiments conducted 20 years later, by Baguhl et al (1995), in which they concluded
that there was evidence in the hyperbolic speed, and the distribution of directions in theimpacts on their detectors, that interstellar dust was flowing into the Solar System.
Moreover Grün et al (1997) were able to estimate a proportion of 30 per cent interstellar
dust at 1.3 AU, and even more at greater solar distances. This section details my simple

167

4.11 A uniform dust flow model

1973 model to simulate of the appearance of an absorption line in sunlight reflectedfrom dust flowing uniformly through the Solar System.

4.11.1 Uniform flow theory
In this analysis:a) The direction of motion of the dust relative to the Sun is assumed to be thenegative of the Sun’s motion towards the Solar Apex. This is equivalent to theassumption that the dust cloud is stationary with respect to the local communityof stars, and the Solar System is moving through it.
b) Focusing effects due to gravity, and scattering by radiation pressure are ignored;
the dust is assumed to be of uniform density and flowing linearly.
c) Scattering is assumed to be isotropic - i.e. diffraction effects are ignored.
The method is to find the shift in wavelength of sunlight reflected from a generalelement in space, and integrate along a line of sight for various elongations. Only the
ecliptic plane is considered.

Figure 4.8 Showing the two Doppler shifts experienced by lightemitted by the Sun, scattered by a dust particle, and observed at the Earth.

Figure 4.8 illustrates the double Doppler shift which occurs in this configuration.

168

4.11 A uniform dust flow model

Referring to the figure, light from the Sun is scattered by the dust particle, and observedat the Earth. On the way, the relative radial velocities of Sun and particle, ΔV1 , andParticle and Earth, ΔV2 both produce Doppler shifts, Δλ1 and Δλ2 . Considering only
components of velocity in the ecliptic plane, radiation emitted by the Sun at wavelength€
λ is ‘seen’ by the particle in its own frame of reference at a wavelength λ + Δλ1 ,€€€where Δλ1 is given by€€

Δλ1 ΔV1=.
λc€where c is the speed of light.

€

€

Δλ1 =

Thus

ΔV1.λ
.
c

An observer on the Earth sees a further shift Δλ2 given by

€
Δλ2
ΔV
= 2 .
λ + Δλ1c €Since Δλ1 is small, we can write λ for λ + Δλ , so

€

λ.ΔV2
Δλ€
.
2 ≈ €
c

€

Hence the total shift in wavelength is€Δλ = Δλ1 + Δλ2 =λ(ΔV1 + ΔV2 ) .
c

We require this as a function of elongation ε from the Sun in terms of the known€ positions and velocities.
ephemeris€169

4.11 A uniform dust flow model

Figure 4.9 Showing the geometry of the scattering, in the plane of the ecliptic only. The diagram showsroughly an October configuration.

In Figure 4.9, which is based on the positions for the October 1st configuration, Vd isthe (apparent) component of the dust particle’s velocity in the Ecliptic plane. Its
direction is 91º ecliptic longitude. Since all speeds are small relative to the speed oflight, the effect of Special Relativity on the angles can be ignored.

λVd cosθ s andcλλΔλ2 = − VE sin ε − Vd cosθ dcc

Δλ1 = −€€€∴ Δλ =λ(−VE sin ε − Vd cosθ d − Vd cosθ s ) .
c

€

170

4.11 A uniform dust flow model

If the solar apex has ecliptic latitude B , and the speed with which the Sun is moving
through the interstellar medium is S , the apparent velocity of the dust Vd is −S cos B ,representing the negative of the Sun’s velocity component in this plane. It has a€numerical value of 11.88 km/s.

€

€

€

It remains to find cosθ d , and cosθ s in terms of ε , the elongation of the dust particle
as viewed from the Earth, and x , the distance of the particle from the Earth.

€

If L is€the ecliptic €longitude of the Sun as seen from the Earth at the time in question,€ of the dust flow has ecliptic longitude 91 degrees:and since the direction€

θ d = 91 − ( L + ε −180)€∴ θ d = 271 − L − ε .

To obtain θ s , referring again to Figure 4.9,€
using the sine rule€giving€Rd=sin ε sin θsinθ =

Rsin εdand, by dropping a perpendicular from the Sun to the line joining the Earth and theparticle,€cosθ =

x − R cos εdThis gives enough information to define θ for all configurations, and θ s is then given€by θ − θ d (the program STREAM used the ARCTAN function with the signs
determined by the two functions above). By this method the Doppler shift for theparticle was determined.

€

€

€

171

4.11 A uniform dust flow model

Figure 4.10 Showing the geometry of the scattering, in the plane of the ecliptic only. The diagramshows roughly an October configuration.

To find the contribution to the line-of-sight integration from an element containing suchparticles, referring to Figure 4.10, the illumination of dust in the element is inversely
2

proportional to d , and if we call the scattering area per unit volume ρ , the lightreflected from the element is proportional to€ρdx × x 2 dΩ
.
d2

€

2

The illumination at the Earth due to this element is inversely proportional to x , so the€eventual light intensity per unit solid angle of view is proportional to

ρdx
.
d2

€

So the intensity dE due to this element, arbitrarily scaled, is given by

€
€

172

4.11 A uniform dust flow model

dE ∝

ρdx
.
x + R − 2Rx cos ε22

We can now determine the ‘average shift” in the usual way as€〈Δλ 〉 =∫ Δλ.dEΔλ∫ dEΔλ
∞

€
Thus

Δλ

∫ x 2 + R 2 − 2Rx cos ε dx

〈Δλ 〉(ε) = x=0∞1∫ x 2 + R 2 − 2Rx cos ε dx

x=0

where Δλ is determined as above.

€
Program STREAM (Appendix 4) uses this method to make predictions for elongations€ at 10 degree intervals. The integral is approximated by dividing x into 1 10 AU steps
and adding up the trapezia, the series being truncated when contributions become lessthan one per cent of the existing total. The output of STREAM was then processed by€€program CONVL (Appendix 5), which convolved the predicted
spectrum with the
instrumental profile of the F-P.
The predicted curves are shown in Figure 4.11, for 1st October and 1st April, along with
the HMR experimental points for each case.

173

4.11 A uniform dust flow model

Figure 4.11 Showing radial velocity predictions as a function of elongation from the Sun for alinear dust flow model, in Spring and Autumn configurations.

174

4.11 A uniform dust flow model

The region around ε = 0 exhibits one of the relatively few examples in Physics of afinite discontinuity. This is to be expected because at low elongations only dust near
the Sun needs to be considered, this dust being so highly illuminated by the Sun, as
dictated €by the inverse square law of illumination. The effect is, however, not a point
discontinuity, because of course the Sun appears not as a point, but with a finite angular
diameter.
For a particle orbiting between us and the Sun, the Doppler shift due to the particle’s
motion with respect to the Sun is in opposite senses on opposite sides of the Sun, as
viewed from the Earth, and as we look at smaller and smaller orbits, the Sun begins tomark a sharp, discontinuous change in Doppler shift.
In fact, the low elongation region is omitted from the diagram for a different reason. If
the particles in this region are small with respect to the wavelength of light, as we havereason to believe they are, diffraction becomes important. Diffraction becomes animportant consideration in general for any particles with small scattering angles – which
occur closest to the Earth along a given line of sight. The half-width of the diffraction
peak is given in radians by:

1.22 λ…. where d is the diameter of the particle.
d
For a 0.3 µ particle and for λ = 0.5µ , this gives€€€

1.22 × 0.5≈ 2 (radians), or about 100 degrees.
€ 0.3

The light diffracted in this case is thus still scattered over a wide range of angles, though€somewhatless than the 360 degrees assumed for isotropic scattering. This, coupledwith the fact that the statistical weighting for these particles is small, means that thisroughly zero-shifted component does not seriously compete with the reflectedcomponents in the integral along a line of sight in the regions we have been able toobserve. Very close to zero elongations this component would eventually become

175

4.11 A uniform dust flow model

important enough to mask the discontinuity around the Sun, by bringing the curves backtowards zero, the effect being more marked if the particles were extremely small.
Perhaps, when spectroscopy of the ZL is performed from a space vehicle, it will beappropriate to return to this discussion. In the meantime there is no need to pay any
further attention to diffraction in any of these models. As a final note on the situation
close to the Sun, there is also reason to believe (Over 1958) that at less than a few solar
radii the particles will be vaporised – another reason to ignore this region for thesecalculations.

4.11.2 Comparison of continuous flow model with experiment
The model gives a good fit to the experimental values for the September-October 1971
period for the high elongations where the rotating model failed to provide a satisfactoryfit. At low elongations in the morning the rotating model gives a better agreement, andat low evening elongations there is little to choose between the two models anyway.
The April 1972 data show some signs of higher shifts at intermediate elongations whichcannot be accounted for by the rotating model, and these could be explained if some ofthe dust were in uniform linear flow; it would make sense that in these elongations weare looking at a much reduced density of orbiting particles, as assumed in all themodels, and a constant density of interstellar immigrants, giving them a relatively
higher percentage of the total in this region. Again, the rotating model satisfactorily
predicts the low elongations. The assumption of the presence of a greater percentage offlowing dust in September-October 1971 than in April 1972 could account for the lackof consistency in fit to the rotating dust model. The Autumn continuous flow curve in
Figure 4.11 has another interesting feature: apart from giving a creditable visual fit to
the HMR data, it gives the right kind of morning-evening asymmetry, and a crossoverwell into the morning side, giving a recessional speed for the Gegenschein. Thesefeatures were both to be confirmed later by East and Reay (1984), as will be discussed
in section 4.13.2.
It seems, then, that a combination of the rotating and continuous flow models might
give the best fit to these data. A ten per cent contribution of continuously flowing dust
produces an effect similar to that seen in our results, but higher quality data, over alonger period, would be needed, to put any certainty into this estimate.

176

4.12 Contribution from sources distant from the Solar System

4.12 Estimation of the effect of a contribution to the Zodiacal
Light from sources distant from the Solar System
4.12.1 Computation
Supposing that the continuum spectrum from background starlight contains some kindof absorption feature at 5183.6 Å, it is possible to estimate the ‘shift versus elongation’curve which would arise from such radiation as observed from the Earth. The shift ofthe line, corresponding to a relative radial velocity, depends on the combination of theSun’s proper motion in the local stellar group, and the Earth’s orbital motion. Since all
the emission takes place at a large distance, no line-of-sight integration is needed, and afurther simplification is introduced if we are interested only in the plane of the ecliptic.
Only components of velocity in the ecliptic plane are considered here.
We require the apparent radial velocity of distant matter as a function of elongation,
which is simply the negative of the ecliptic plane component of the Earth’s total motionin the direction specified by this elongation.
The geometry is a simpler version of the continuous flow model, and is shown in Figure
4.12. If the Sun’s proper motion is S km/s and the direction of the apex of this motion
has ecliptic latitude B , the numeric velocity component in the ecliptic plane is S cos B ,directed towards ecliptic longitude A . Coordinate conversions from Right Ascension€ longitude and latitude were made using tables. For any day
and Declination to ecliptic€ the ecliptic longitude of the Sun, C degrees, is given €in the year
in the Astronomical€ latitude is, of course, zero. From this, the direction of theEphemeris. The ecliptic
Earth’s motion, at speed E km/s relative to the Sun, is C − 90 (degrees ecliptic€ as vectors in Figure 4.12 in which the angles A
longitude). These velocities are shown
and C are measured relative to 0 degrees ecliptic longitude.

€

€

177

4.12 Contribution from sources distant from the Solar System

Figure 4.12 Showing the Doppler shift from light emitted by distant sources,
observed from the Earth, in the ecliptic plane only.

The angle, α degrees, between the two components of velocity is

α = (C − 90) − A

€

i.e.

α = C − A − 90 .

€ the cosine rule for the component triangle, since its obtuse angle is 180- α , theUsing€
R , of the component in the ecliptic plane of the Earth’s total velocity, ismagnitude
given by
2

2

2

€1

R = E + S cos B + 2ES cos B cos(C − A − 90) 2{€This is directed towards ecliptic longitude L where€E
L = A − sin−1€ sin(C − A − 90) ,R
obtained using the cosine rule in the same triangle.

€

}

178

4.12 Contribution from sources distant from the Solar System

The elongation of this vector from the Sun

εR = L − C

ε R = A − C − sin−1

and soEsin(C − A − 90) .
R€
Values of E and S were taken from Astrophysical Quantities and R and ε R were€

calculated for October and April. The apparent component of velocity of distant matterat an elongation ε from the Sun is evidently

€

€

€

€

−R cos(ε − ε R ) .

€The wavelength shift δλ for a given radial velocity is found by the usual formula for€

optical Doppler shift at a wavelength λ

€

δλ v
=
λ€ c

where v is the radial velocity and c is the speed of light. Thus we have€

δλdis tan t matter = −€€λR cos(ε − ε R ) .
c

This function is shown plotted in Figure 4.13 for the two periods of observation. The€ figures are
Time

October 30th
April 15th€− λR c (Å)

εR

0.3
0.6

73° W

€

97° W

The Earth’s velocity – approximately 30 km/s – is the dominant factor in the Doppler
shift since S cos B is only about 12 km/s. This means that the maximum blue shift isalways on the morning side within about 15 degrees of the 90° elongation position. The

€

179

4.12 Contribution from sources distant from the Solar System

Sun’s velocity adds in the case of the April position, which gives near maximum
amplitude of the curve. October is near the minimum case and shows a small
amplitude.

Figure 4.13 Showing all the wavelength shift data available in 1974,
compared with predicted curves for distant emitting matter.

180

4.12 Contribution from sources distant from the Solar System

4.12.2 Comparison of distant emitting model with the data
Comparison of the curves with our experimental points shows that the rough generalshape of the experimental curves can be produced by such a theory, so it is certainly
worth finding out if galactic starlight does contribute to the spectrum observed, and inwhat proportion.
In the October case, the curve is quite similar to that predicted by the rotating dust cloudmodel but, like the latter, does not satisfactorily fit the data, predicting a cross-overpoint nearly 60° in error, and a maximum blue shift at an elongation where theexperimental points are at around zero shift. The contribution of galactic light in the
October observations must thus be small.
In the April case, this theory predicts very large maxima, not commensurate with mostof the data. A few points are, in fact, found with shifts up to 0.6 Å, and it is possiblethat the spectra for these points were measured when the direction was in line with aregion of strong starlight background. The outskirts of the Milky Way may have been
responsible. Overall, however, the curves fitted to the data do not on the whole agree
well with the predictions of this theory. This, added to the fact that considerable care
was taken to avoid making measurements in areas of the sky contaminated by othersources, leads me to believe that the contribution of a background starlight absorption
line is small enough is most cases to be ignored in our analysis.
This opinion is strongly reinforced by the experimental appearance of line shapes.
These are very near to the theoretical predictions of the program CONVOL, which
simulates the output when a pure ZL spectrum is input to the system. Diffuse starlightis by no means entirely composed of light from stars with pronounced Fraunhoferfeatures in their spectra. Even if it were, in a particular direction, the lines should besmeared by the combination of proper motions; and so an absorption line, if it exists at
all in the spectrum of distant aggregate starlight, would be expected to be shallow and
wide. The very fact that our observed shapes are not significantly shallower than thepredictions from local dust particle models suggests again that backgroundcontributions are, on the whole, small. It is worth noting here that for the future it
would be useful to confirm this view of the spectral content of distant starlight, both in181

4.12 Contribution from sources distant from the Solar Systemtheory, in different populations of stars and regions of nebulosity, and in practice, bydirect spectroscopy of such regions, without the interference of the ZL. Of course, thisrequires a viewpoint outside the Solar System.

4.13 Further analysis of the data by comparison with newradial velocity data and model theory available in 2007
As discussed in Chapter 1, there have been long periods of inactivity in Doppler studiesof the ZL since 1974. Nevertheless, some interesting results have now been added tothe data in this field, and it is very informative to study these for the first time in relation
to the complete HMR observations. I present them here in chronological order.

4.13.1 Observations of Fried 1977

Figure 4.14 Comparison of our data with that of Fried 1977.

182

4.13 Further analysis of the data

Figure 4.14 shows the results of Fried (1977). It is immediately apparent that there islittle or no agreement here between the two sets of measurements. Reference to Figure
4.5 suggests that Reay’s model with p = 2 is a fairly good fit to Fried’s points, which
paints a picture of large numbers of large dust particles rotating in prograde orbits, atspeeds close to the orbital speed of the Earth in the case of particles at around 1 AU,€
giving those very low wavelength shifts around 90 degrees elongation. Indeed Fried
himself deduced an ensemble of dust particles in hyperbolic orbits, with a size
dependence on orbital radius r given byn(r ) = r −1.3, and a corresponding dependence of velocity€

V (r ) = r −1.5 .
€If these data had come to my attention in 1977, I would doubtless have felt depressed,€ since they do not in any way support our results. It is therefore something of a comfort
that as time went on, no other observers’ findings conformed with these of Fried. This
is not, of course, to say that there is anything wrong with these observations. There is
still, as we have discussed, debate about the variability of the ZL, and it may well bethat, since no-one else was observing its Doppler signature in 1977, Fried saw atemporal effect.

4.13.2

East and Reay (1984)

Figure 4.15 shows the points published by East and Reay (1984) on the single
wavelength shift versus elongation graph, plotted in direct comparison with the HMR
data.
These data have an improved signal-to-noise over ours – the product of a more stable,
servo-controlled Fabry-Perot etalon, and a greater number of observations to average.
Overall they show a remarkable congruency with the HMR data, with the qualifications
that the wavelength shifts are slightly smaller in the range 60º East to 60º West, and themorning-evening asymmetry is even more pronounced. East and Reay, from apolynomial curve fitted to these results, reported a confirmation of a predominantly
prograde orbiting Zodiacal Cloud, with no evidence of retrograde component, in183

4.13 Further analysis of the data

elliptical orbits, rather than hyperbolic, and a recession of the Gegenschein of about 2.5
km/s. The Gegenschein results were supported by an extra body of data specificallyamassed to reduce errors.

Figure 4.15 Wavelength shift vs. elongation: the results of East and Reay (1984),
compared with the Hicks, May and Reay data.

4.13.3 The WHAM observations, and theoretical model predictions
In 2004, Reynolds, Madsen and Moseley published some new data, collected using the
WHAM spectrometer, designed for H Alpha mapping, at Kitt Peak, Arizona. The
spectrometer itself was a dual Fabry-Perot, 15 cm in diameter, of superior resolution(0.2 Å) and high throughput. The results were published in Reynolds et al (2004), andanalysed with reference to models by Ipatov et al (2006) in Madsen et al (2006). A
great improvement of signal-to-noise ratio over our measurements was achieved byfeeding the spectrometer from a 0.6 m telescope, the whole assembly being remotely
controlled from Madison, Wisconsin.

184

4.13 Further analysis of the data

Figure 4.16 Wavelength shift data and theoretical curves from Madsen et al (2006),
plotted in comparison with HMR data.

The observations were obtained quickly, on just two nights in November 2002.
Because of the higher resolution, the WHAM team were able to identify, and subtract,
non-MgI emission lines from the atmosphere, contaminating the MgI absorption line,
again improving the reliability of the data. It is possible that these unseen emissions
had an influence on our data points – perhaps even producing the morning asymmetry
and the red shift of the Gegenschein, though this would be hard to prove. The WHAM
results, shown in Figure 4.16, are shown plotted alongside the HMR data. The generalagreement between the two sets of observations is quite good, although there is atendency for the HMR data to present higher shift at intermediate longitudes, and the
WHAM data does not show any morning-evening asymmetry, or any shift at 180º for
the Gegenschein.

185

4.13 Further analysis of the data

In 2006, Ipatov et al published their model of a Zodiacal Cloud, starting from theerosion of comets and asteroids, as mentioned in Chapter 1, and developing orbital
velocity elements by simulating their dynamical behaviour with time. Some of theresulting curves are shown in Figure 4.16, and it can be seen that they fit the WHAM
observations well. I decided to see how the HMR radial velocity data would look, ascompared to the Madsen et al (2006) curves, if they were binned into larger groups of
averaged scans (next section).
We have already seen the WHAM line-width data in Figure 4.4, and remarked that there
was some similarity with HMR in the grouping of the points.

Certainly the

disagreement in the wavelength shift data between HMR and WHAM is not very big,perhaps within the probable errors of the HMR data; but reconciling WHAM with theEast and Reay data, with its smaller error bars, is not so easy (Figure 4.17).

Figure 4.17 Comparison of East and Reay (1984) with WHAM data.

186

4.13 Further analysis of the data

The evidence is clearly conflicting here; East and Reay’s data shows a distinct
divergence from WHAM at all elongations which is not inside the predictions ofprobable error. Since our HMR data conforms more closely to East and Reay than
WHAM, there is still evidently a need for a new set of data, to determine if thedifferences are a conflict, or the result of seasonal or other variations. The East andReay results, in almost every case, show a shift towards the red as compared with
Madsen et al (2006). As we saw earlier, the only model which has a chance of
explaining such asymmetries between morning and evening, or Spring to Autumn, is
my continuous flow, interstellar dust model. Could it be that we are seeing a
combination of rotating and free-flowing dust in the ZL? If the difference between
Spring and Autumn radial velocities suggested by our data is real, it is possible that the
WHAM experiment, a snapshot in time at the November configuration, was simply
fortunate in catching the radial velocities in a neat, symmetrical pattern, perhaps in aperiod of low density in the interstellar dust component. Perhaps only time and still
higher resolution and signal-to-noise will tell, but there would clearly be good reason to
try to monitor the radial velocity versus elongation curve all the year round. There isone more piece of reductive analysis on the HMR data which can be done: binning.

4.14 ‘Binning’ of the HMR data
Bins of 20-degree spans in wavelength were chosen, and the mean shifts in each 20degree interval were calculated. The results are shown in Figure 4.18.
This figure shows only elongations from zero to 180º, because the plotted points forevening and morning have been ‘folded’ together, reversing the sign of the evening
observations, as many previous authors have done, assuming the shifts to besymmetrical about the antisolar point. The ‘ghost’ points are the original HMR data.
Three sets of binning averages are shown, the results for 1971 only, for 1972 only, andfor both sets averaged together. Except for the region 110º-130º, the HMR observations
consistently show a shift about 5 km/s larger than the Madsen/Ipatov curves. Again thefit seems quite good for a Reay curve in which p = 5 . The separate binned HMR
points underline the view that there is a real difference between morning and evening
radial velocities, and a time dependence.

€

187

4.14 'Binning' of the HMR data

Figure 4.18 The result of binning the HMR data in wavelength intervals of 20º.

Finally, Figure 4.19 shows a direct comparison between the consolidated HMR data andthe WHAM data, both folded about the antisolar point. The folding reveals that there
actually is a morning-evening asymmetry in the WHAM points – not immediately
apparent in the un-folded figures. The morning points seem to cross the zero line atabout 30º elongation, and the evening points at about 50º – a discrepancy well outside
the limits of the predicted uncertainties in the data. The HMR data shows a crossing at
about 40º, but is lacking in enough low elongation data to be sure. Elsewhere the188

4.14 'Binning' of the HMR dataagreement between the two sets of data is good except between 60º and 110º - where
our radial velocities are almost twice that indicated by the WHAM data. It has to beemphasised that all the information on time variation, if there is any, has been
eliminated from this graph.

Figure 4.19 Binned folded HMR data, averaged in wavelength intervals of 20º,compared with folded WHAM data.

4.15 Variations with time
I did not expect to find evidence of short-term changes in the parameters of the Zodiacal
Cloud. But the differences between our two periods of observation (at opposite ends ofthe Earth’s orbit) are so pronounced that it is plainly worth checking in the future to see

189

4.15 Variations with time

if the differences can be reproduced. Throughout the history of the observation of the
ZL there have been reports of seasonal and other variations in intensity of both morningand evening ZL. But there has not been a consensus on what these variations are, or themagnitude of the changes, and since very few truly long-term projects have been
undertaken, it has not been established whether the fluctuations are cyclic or not withthe season, or due to unique events, such as solar flares or visits from comets, or indeed,as discussed later in this Chapter, due to influences beyond the Solar System.
Variations in intensity were recorded by naked-eye observations as long ago as 1856, by
G. Jones in his voyage from the USA to Japan and back on a boat, which I mentioned in
Chapter 1 (Figure 1.1). Jones’s meticulously drawn nightly contour maps of theZodiacal Band show what he thought was a dependence of the ecliptic latitude of thecones on the latitude of the observer. Jones interpreted this dependence as an indicationthat most of the Zodiacal Cloud was attendant on the Earth, rather than the Sun.
Though this became an unfashionable view in the second half of the 20th Century, theexistence of a component of ‘local’ light-scattering dust can still not be ruled out,
especially since the discovery (Reach et al 1995) of the ‘Earth Tail’ of dust,
gravitationally channelled by the passage of the Earth on its journey around the Sun. It
is also possible, however, that the variation Jones saw was not latitudinal, but seasonal,
and this might turn out to be the earliest observed manifestation of the departure of the
ZC from the invariable plane, which has been documented in recent satellite
experiments (Hauser et al 1984).
E. O. Hulbert (1930), in a much later analysis of the Jones observations, wrote that
Jones had witnessed 23 unusually bright periods of the ZL, 16 of which took place 3days after a magnetic storm. Hulbert also noted other enhancements preceded bymagnetic activity. He notes, however, that some of the variations might be interpreted
as overall sky brightness changes, as a true separation of ZL component does not seemto have been made. It seems likely to me that Jones himself would have placed moresignificance in his observations of the variations shape of the ZL than those of overall
intensity.
In more recent times, seasonal oscillations in brightness, of magnitude about 20 per
cent, were reported by Elvey and Roach (1937) using photoelectric measurements.

190

4.15 Variations with time

In the evening ZL they found a maximum in April-May, and a minimum in JanuaryFebruary; and an opposite effect in the morning ZL. These conclusions were supportedby their reduction of visual Japanese observers.
Divari (1955), found seasonal variations in which the morning ZL intensity increased in
September to November, and the evening intensity decreased from January to March.
These results clearly do not concur with Elvey and Roach’s. As noted in Chapter 1,
Sanchez (1974) found a correlation between ZL brightness and solar activity.
The question of variability of the ZL seems central to interpretation of the Doppler data
so far gathered in 2007, to account for the discrepancies revealed in this thesis between
all currently available sets of results.

4.16 Summary of the implications of the 1970-4 work,and pointers for the future
It is clear that a reasonable fit to our data can be obtained using a model of the Zodiacal
Cloud comprising dust particles, some in orbit around the Sun, and some drifting
through the Solar System, possibly to be identified with interstellar dust.

The

differences between the two periods of observations, if they are real, may be attributed
either to a seasonal effect, corresponding to an asymmetry of the dust cloud with respect
to the Earth’s orbit, or to short-term fluctuations of the orbiting or drifting component.
Such short term variations in the former might be due to the visits of comets, orvariations in Solar Wind. In the latter, the differential would indicate either that there
are short-range inhomogeneities in the interstellar medium, or that, in its passage
through the Solar System, the flow of such material is disrupted, for instance by theplanets themselves.
In future work the data from high elongations may be the most useful, since a purerotating model can account for the shape in all cases for which ε ≥ 70 (degrees ofelongation from the Sun). Our results are consistent with the theory that a higher
percentage of drifting dust was present in September-October 1971 than in the April
period the following year.

€

191

4.16 Summary of the implications of the 1970-4 work

More data is still required, critical for distinguishing between models, and hence forbelief in the conclusions we are drawing, and more precise determinations of the actual
spectrum profiles, rather than simply averaged shifts, will in the future yield better
mapping of the motions of the interplanetary dust cloud.

192

5.1 Current Zodiacal Light Research

Chapter 5
Current developments and future plans
In this chapter, I will briefly review the experiments currently planned, or inoperation, in Solar System Astronomy: rocket-borne, space-borne, and groundbased observations, which are likely to throw new light on the structure andnature of the Sun’s dust cloud. I will also touch on present unresolved issues which
warrant attention, and sketch plans for future observations which Dr. Garik Israelian
and I plan to make, to extend and refine the Doppler radial velocity measurement
technique. We hope to map dust kinematics to add new detail to the latest picture of theZodiacal Cloud, as revealed by IRAS, COBE and ISO.

5.1

Current Zodiacal Light research

CIBER
The Cosmic Infrared Background Experiment (CIBER) (Bock et al, 2006), is a rocketborne instrument designed to search for signatures of primordial galaxy and starformation in the cosmic near-infrared extra-galactic background, in the wavelength
region 0.8 to 2 microns. Its launch, funded by NASA, is scheduled for September 2007.
It is hoped that CIBER will settle the question of whether the DIRB really is the echo of
the first stars and galaxies. CIBER consists of a wide-field two-colour camera, a lowresolution absolute spectrometer, and a high resolution narrow band imaging
spectrometer. It will trace the Zodiacal Light with the low resolution spectrometer,using the intensity of scattered Fraunhofer lines, to provide an independentmeasurement of Zodiacal emission, and a new check on the Kelsall et al Zodiacal dust
models, which were used to compare with the results of the Diffuse Infrared
Background Experiment (DIRBE) of the Cosmic Background Explorer (COBE). It isalways necessary to subtract an accurate estimate of the intensity of ZL emission inorder to map the Galactic contribution and the level of emission from elsewhere – theInfrared Background. Models of the levels of ZL contribution are still so uncertain that
estimates of the IRB are, at present, tentative. Being essentially a low resolution

193

5.1 Current Zodiacal Light Researchexperiment, CIBER will not be able to do spectroscopy, as such, of the ZL. But the
CIBER team hope to accurately assess the contribution of local (z = 1-3) galaxies, andhence to determine the excess unaccounted for. In this residual radiation, CIBER will
be searching for the newly-predicted signature of high luminosity primordial stars andmini-quasars - a red-shifted Lyman cut-off feature between 0.8 and 2 microns.
PLANCK
It will be interesting to see what view the Planck experiment, due for launch in 2008,
will have of the Zodiacal Light. The satellite will be put into a Lissajous orbit aroundthe Sun-Earth libration point L2 , approximately 1 AU from the Earth, which will give ita significantly different perspective on the Zodiacal Cloud.

And, ‘seeing’ atwavelengths over 300 microns, it is possible that the ZL will look very different in€ construct all-sky maps in nine frequency bands, covering the range
shape. Planck will
30 GHz to 857 GHz, with relatively high sensitivity and resolution (Maris et al, 2006).
According to the Planck ‘Bluebook’ (ESA, 2005), Planck is expected to detect Zodiacal
Light at 857 GHz and possibly 545 GHzd; at these wavelengths the Zodiacal emissionis, until now, poorly constrained by observations, and its level uncertain by more than afactor of two. At lower frequencies than this, levels of emission from the ZL are
expected to be very low. It is hoped that accurate mapping of the dust emission in thiswavelength region will give information on the geometrical distribution of the largest(up to 1 cm) grains in the dust cloud, in three dimensions, and as a function of season.
The detection of thermal emission at wavelengths as long as this makes Planck sensitive
to dust colder than has been previously mapped; finding such emission could imply theexistence of an extension of the Zodiacal Dust Cloud into the outer reaches of the Solar
System, perhaps into the Kuiper Belt. Of course, even if the cloud only extends to theasteroid belt, there would be Rayleigh-Jeans radiation from warmer dust at 350 microns.
The signature of colder, more distant dust would be an enhanced intensity ofsubmillimetre emission. The possibility of Planck detecting the IRAS dust bands does
not appear to have been discussed; however it is hoped that it will be able to map thecometary dust trails reported by IRAS.

194

5.1 Current Zodiacal Light Research

SPITZERThe Spitzer orbiting infrared telescope was named after the first astronomer (Lyman
Spitzer in 1946) to suggest putting a telescope in space. It has a sensitivity 1000 times
better than ISO, and has full wavelength coverage in the mid-IR (Werner et al, 2004).
The “First Look Survey” of Spitzer (FLS) has investigated the mid-IR Backgroundbrightness in the ecliptic plane. The sensitivity of Spitzer is background-limited, anddominated by Zodiacal Light at all but the longest wavelengths and lowest Galactic
latitudes. Spitzer carries a short wavelength Infrared Array Camera (IRAC), which
provides absolute photometry, relative to the brightness of a dark sky patch, at a rangeof four wavelengths spanning 3.6 to 8 microns. This instrument is sensitive to emission
from evolved galactic systems, and has produced many beautiful multi-‘colour’ imagesof many nearby galaxies.

The 24 micron long-wavelength Multiband Imaging

Photometer for Spitzer (MIPS) has provided absolute photometry of ZL emission. The
measured ZL background was found to be within 5 per cent of Spitzer model
predictions at 24 microns derived from COBE/DIRBE data. The total backgroundbrightness at 8 microns was 36 per cent brighter than the model predictions. However,this is thought to be consistent with known intrinsic variations of the ZL as a function ofviewing geometry (Meadows et al 1997).
Recent reports of analysis of Spitzer data (Reach 2005) indicate that it is indeed
achieving success in the pursuit of refinement in the known structure of the Zodiacal
Cloud. They say, “Using the new infrared spectral and imaging capabilities of theSpitzer space telescope, we find dust trails from short-period comets, new structures inthe zodiacal light, weak 9-11 micron silicate spectral features from some comets andthe zodiacal light, and the color temperature of cometary nuclei, comae, and dust trails.
The high frequency of dust trail detections for short-period comets suggests they are ageneric phenomenon and that most such comets lose mass primarily in the form of
meteor-sized particles. The color temperature of the trails is higher than a grey-body,while their dynamics suggest they must be of order mm-sized, so they must maintaintemperature gradients across their small surfaces. The zodiacal light silicate feature
requires a contribution from small amorphous silicates, though the continuum requires
larger particles.” This quote has direct relevance to the work I hope to undertake in the
coming year.

195

5.2

5.2 Objectives of our proposed new work in Zodiacal Light

Objectives of our proposed new work in Zodiacal Light

5.2.1 Doppler measurements on scattered light at optical wavelengths
Dr. Garik Israelian (of the Institute of Astrophysics, La Laguna, Tenerife) and I believe
it would be very instructive to look at Doppler shifts in regions away from the ecliptic
plane, and see if we can detect a higher proportion of retrograde material. We haveseen that high radial velocities might also be an indication of a significant Interstellar
component. If we can obtain good enough quality spectra in a slice perpendicular to theecliptic, including the ecliptic pole and the Gegenschein, it should be possible to choose
between these hypotheses.
We have seen that the IRAS observations, interpreted by Low et al (1984), indicated theexistence of dust bands, producing ‘shoulders’ superposed on the smooth photometricshape of the traditionally known Zodiacal Light. We propose to look for the dust
bands’ dynamic signature, by measuring Doppler shifts in a slice of the ZL across the
ecliptic plane, at an intermediate elongation, again using a ground-based optical andmid-infrared spectrometer. A sudden change in radial velocity at some point in thisslice would be confirmation of the existence of such a discrete component in the cloud,and would offer the first opportunity to determine its motion; this could then becompared with the motion of the putative parent asteroidal families.
As for instrumentation, with the benefit of 30 years’ hindsight, it seems to me that, as
far as maximising flux is concerned, we, that is, all the F-P radial velocity observers of
the time, missed the obvious. The Fabry-Perot instrument was used because of its highthroughput at a relatively high resolution, but, by using the pre-monochromator, we
actually threw away 3,000 Å of information in the visible alone, all of which contained
absorption lines shifted by the same amount, and all of which therefore potentiallycould have given us more information on radial velocity. In admitting only actually 6 Å
spectral range, we rejected 99.8 per cent of the available useful signal. In addition, weused only a small coelostat to collect light. The justification was that we wanted to lookat a large area on the sky to achieve large flux-levels, and it would have been hard tofind a telescope with a comparable field of view, of one degree, on the sky. But in factlarger flux levels can be obtained with large mirrors with a lesser field of view, more

196

5.2 Objectives of our proposed new work in Zodiacal Light

than compensating for this gain. In future I would like to correlate the modification of
the whole spectral range available, and thus vastly improve the signal-to–noise in theobserved spectra. Moreover, with modern echelle technology, there is every reason toemploy the largest collector of light available to feed the spectrometer in future. These
two improvements ought to allow not just comparison of an ‘average shift’ with theory,
but detailed comparison of multiple spectral line shapes with theoretical models,yielding a much more thorough and accurate mapping of radial velocities. Exploratory
experiments will be made in July 2007 with an echelle spectrometer attached to a 3.5 m
telescope to test the validity of these arguments and see what quality of results can beobtained.
Recently, Bernstein, Freedman and Madore (2002) detected extragalactic backgroundlight using an echelle spectrograph at the 2.5 m Du Pont telescope in Las Campanas(Chile). These authors have developed a technique that uses the strength of theZodiacal Fraunhofer lines to identify the absolute flux of the Zodiacal Light in thecomplex night sky spectrum. The analysis is much simpler if we are only interested in
the Doppler shifts of Zodiacal Fraunhofer lines.
Dr. Israelian and myself are planning observations with the SARG spectrograph locatedat the Nasmyth focus of the 3.5 m Galileo telescope in La Palma, Canary Islands.
SARG is a high efficiency echelle spectrograph designed for the spectral range 370 to
900 nm, and for resolution from R=29,000 up to R=164,000. Both single object andlong slit (up to 30 arcsec) observing modes are possible. A dioptric camera images the
cross-dispersed spectra onto a mosaic of two 2048 by 4096 CCDs. We plan to use along slit mode (aperture 5, providing an area 1"07 by 26"7 arcsec) which provides a
resolving power of 43,000.
To minimize read-out noise, we will bin the data on the chip in intervals of 4 pixels inthe spatial direction, and average over the full extent of the slit in the data reduction. We
will use 30 minute exposures which are expected (given the results of Bernstein et al
2002 and the Exposure Time Calculator of SARG) to provide a signal-to-noise ratio
between 60 and 80 per resolution element. We shall use a cross disperser CD1 (600g/mm, 4840 angstrom, order separation 10.7 arcsec) which provides spectra in the range
3600-5140 A in 50 orders. This spectral window is not severely contaminated by sky

197

5.2 Objectives of our proposed new work in Zodiacal Light

glow (Bernstein et al 2002), being in the blue part of the spectrum, and we hope to beable to observe ZL lines directly in these spectra without insurmountable problems insky subtraction.
The solar Fraunhofer lines are clearly visible in this part of the spectrum and they will
be Doppler-shifted due to the scattering by dust particles. We will use the Solar High
Resolution Spectral Atlas of Kurucz et al to identify and study the Fraunhofer lines.
We are not interested in absolute flux determinations. This experiment will allow us toincrease the precision of Doppler measurements due to the fact that we will beobserving many Fraunhofer lines simultaneously. We aim to observe the ZL at various
latitudes and elongations, and will study the velocity fields of the dust bands in visible
light. Our study will allow us to constrain models of the ZC, and dust band formation.

5.2.2 Doppler shifts of emission lines in the infrared
Investigating the ISOCAM spectra, Reach et al (1997) found a hint of a 9-11 micronfeature, and proposed that the particles producing the ZL are composed of silicatessimilar to those found in the tail of P/Halley. The mid-infrared spectra of many comets
show structured 9-11 micron emission bands, and the spectrum of the dust cloud around
Beta Pictoris has been found to exhibit a comparable structure. The shape of the ZL
silicate emission line appears to be similar to those of comet Halley and thecircumstellar dust around Beta Pictoris, but the ZL silicate feature is superposed on arelatively brighter continuum. It is very important to confirm the detection of thesilicate line at 9-11 microns. Further work with Spitzer may reveal more structure inthis feature. We propose to look at it from the ground.
We plan observations of the silicate feature with VISIR on the 8.2 m VLT (UT3) inChile. Using a high resolution mode (providing resolution 25,000 at 10 microns), we
will try to observe in the 8-12 micron atmospheric window. We propose to use a longslit mode: a slit width of 1 arcsec, and slit length 32.3 arcsec. The sky (airglow) lines
will be subtracted by the "dithering" and "nodding" method (see for example, theGemini Near Infrared Spectrograph website 2007).

198

5.2 Objectives of our proposed new work in Zodiacal Light

These observations will allow us:1) to verify the presence of the silicate emission feature.
2) to study the line profile and its variability.
3) to study the silicate feature in ZL dust bands.
4) for the first time, to attempt Doppler shift measurements on an emission line, if there
is enough fine structure in the feature to allow this.
We feel that, in the future, the kinematic Doppler-shift signatures of moving dust
particles will provide a complimentary source of information, alongside continuing
photometry, on the structure of the Zodiacal Cloud: this will lead to more confidence
not only in our understanding of interplanetary dust and its origins, but also in ourassessment of the emissions from beyond – from Interstellar, Intergalactic, Cosmic, andBackground sources.
Almost 35 years after my first experiences with a Fabry-Perot spectrometer aimed at theZodiacal Light, there are still excellent reasons to continue, with ever improving
techniques, to map Radial Velocities in the Zodiacal Dust Cloud.

5.3

Concluding remarks

5.3.1 Summary
The results of my work in 1970-4 were, as planned, the first comprehensive survey of
Doppler shifts in the Zodiacal Light, covering elongations in the Ecliptic plane from 25º
East to 25º West. The spectra and plots of wavelength shift versus elongation stand as abody of work which continues to be cited by those who have followed in the field.
There is fair agreement between my results and those of East and Reay, and Reynolds etal, who both worked later with higher power equipment. The differences between thebodies of data still pose questions which further observations will be able to clear up:
the question of whether there are variations over time in the kinematic structure of the
ZL and if so, whether they are seasonal, and apparent, caused by the Earth’s journeyaround the Sun, perhaps through a continuous flow interstellar component of dust, or if

199

Concluding Remarks

they reflect actual structural changes in the Zodiacal Dust Cloud, perhaps due toasteroidal events or visits from comets.
My conclusions at the time were that an emission component at MgI existed, generated
in the Earth’s atmosphere, that the Zodiacal Cloud consisted of dust particles of theorder of 1 micron, in orbit around the Sun, mainly in a prograde sense, and that themorning-evening asymmetry we found might be due to components of dust NOT in
orbit around the Sun. We found no evidence for contra-rotating dust, but at theresolution used I do not consider that the presence of a small percentage of such dust is
ruled out. The recession of the dust in the area of the Gegenschein was, and is, anunsolved mystery. This effect was confirmed by East and Reay, but not seen at all byReynolds et al. So it is not entirely clear whether it really exists or not, or if it
sometimes exists! The results of Fried are not congruent with any of the other three setsof data, but apart from these, the rest of my, and our, conclusions of 1974 remain
supported by all other subsequent experiments.

5.3.2 Visions
Often it is hard, among the wealth of technical details published on such a subject, toform, in the head, a clear picture of the physical shape of the Zodiacal Cloud – the kindof graphic visual we would like to present to our children, and those who follow us.
Figure 5.1, from Stark (2006), is a rare attempt in the published literature to provide ageometrical visualisation of the structure of the dust cloud.

The figure isdiagrammatical, but clearly illustrates the three dust bands, showing them as flat rings,
and makes no attempt to show the ‘hollow torus’ indicated by evolutionary models.
It is highly likely that this representation will have to be updated, as new and bettermeasurements are achieved.

200

Concluding Remarks

Figure 5.1 From Stark (2006). A geometrical representation of the 1998 DIRBE model of theZodiacal Cloud. The red spheroid represents the smooth diffuse component (no structure attempted);
the blue rings are the three (presumed asteroidal) dust bands; the green ring is the circumsolar
Earth-resonant ring, showing a ‘Blob’ following the Earth (coloured blue).

The final three figures which follow are ‘artist’s impressions’.

Figure 5.2 is a

representation of how the consensus of Zodiacalists pictured the Sun’s dust cloud in1970 – quite smooth and diffuse, roughly lenticular in shape, symmetrical about theecliptic plane, its density decreasing quite quickly beyond the orbit of Mars. Figures
5.2 and 5.3 are speculative visualisations of the information currently available in 2007.
In Figure 5.3, the various currently accredited components of the dust cloud areillustrated: again a relatively smooth continuous ellipsoidal (lenticular) main cloud,perhaps spreading a little less in the direction perpendicular to the ecliptic than the Fan
Model suggests, but now showing significant structure. Each inner planet has afocussing resonant effect, producing a system of circum-solar rings at the radii of theplanets - and corresponding depleted areas between the planetary orbits. The innermost
rocky planets also produce 'blobs', similar to, but smaller than the blob which has been
observed following the Earth, and they have an overall effect, over time, of producing'lumpiness' in the structure (not yet observed). But most spectacularly, the smooth dust
cloud is now flanked by discrete inclined dust bands, originating from collisions in theasteroid belt. This visualisation (Figure 5.3) follows the interpretation of Stark (2006)

201

Concluding Remarks

in Figure 5.1, the IRAS observed photometric irregularities each being attributed to a
separate narrow dust band. Some asteroids can be seen, including some of the families
which have produced the dust bands, outside the orbit of Mars. There is an additional
faint huge, cool dust ring, out towards the Kuiper belt, again, not yet seen! The angles
of inclination of the rings are depicted schematically only.
Figure 5.3 shows an alternative and probably more plausible interpretation of the IRAS
dust band discoveries. Models developed by Dermott et al (private communication)
show how a single collision between large asteroids can initially produce a band of dust
debris distributed along the path of the ‘parent’ bodies, similar to those shown in Figure5.3. But over a few million years the models show the individual orbits of the particles
precessing at various rates, until eventually the major axes of their elliptical orbits are
spread randomly around the Sun, producing a smooth, fat toroidal distribution. The
orbits of these particles may be thought of as a component in the ecliptic plane, plus a
synchronous sinusoidal “North-south” movement. Hence, viewed from the Earth,which is situated inside the torus, the dust particles spend more time at the edges of the‘doughnut’ than inside, so we see peaks in optical depth at the top and bottom edges of
the torus. Evidently these peaks would be symmetrically placed above and below the
ecliptic, which, in the case of the 10º IRAS bands, is certainly what is observed. Figure
5.4 shows one such torus, a highly evolved remnant of an asteroidal collision. The
complete picture may require the addition of at least two more toroidal components.
Perhaps some combination of Figures 5.3 and 5.4 would truly represent the appearanceof the complex circum-solar dust cloud that will be eventually photographed from aremote space craft in the future. In addition to the asteroidal debris, a comet is seen
approaching the Sun, a reminder that quite possibly the major fraction of the dust isactually due to comets. Indeed, to this picture should probably be added a whole fine
tracery of cometary dust trails, the ejecta which follow the orbital paths of the cometson their journeys to and from the Sun. And, just possibly, we should add a fairly
uniform ‘background’ of dust flowing through the Solar System, a consequence of theSun’s journey through interstellar space; one wonders how much of a contribution thismight be making to current estimates of the Diffuse Infrared Background. If we could
accurately map Doppler shifts in this residual drifting component, we would know.

202

Concluding Remarks

Perhaps in the near future, kinematic studies of the material around us will resume theirrightful place alongside photometry, in unravelling the secrets of the formation andevolution of the Zodiacal Dust Cloud.
The ‘artist’, by the way, is me.
Brian May, 2007.

203

Concluding Remarks

204

Concluding Remarks

205

Concluding Remarks

206

Concluding Remarks"In the sweat of thy face shalt thou eat bread, til thou return unto the ground;for out of it wast thou taken:
for dust thou art,
and unto dust shalt thou return."from the Bible, Genesis 3:19 (King James Version).

207

Computer programs

Appendix 1 Program ZLPLTProgram to process a night’s raw data, subtracting the dark count, dividing out the
premonochromating filter, dividing by the compensation channel count, and outputting
the fully reduced scans on to cards.

208

Computer programs

209

Computer programs

210

Computer programs

Appendix 2 Programs which fit curves to the fully
reduced data points
LSQFIT accepts the punched card fully reduced data and calculates a least-squares
fitted polynomial curve of any specified order (to a maximum of six, in this case).
Output is on cards, to be read as input by the programs in Appendix 2, which plot the
curves and estimate the mean wavelength shift.
GAUSSN accepts the same data cards and calculates a least-squares fitted Gaussian
curve, outputting parameters: mean wavelength shift, line depth and line width. Output
is on cards, which are read by the program GAUSPT, a plotting program.

211

Computer programs

212

Computer programs

213

Computer programs

214

Computer programs

215

Computer programs

216

Computer programs

217

Computer programs

218

Computer programs

219

Computer programs

220

Computer programs

Appendix 3 Programs which draw the fitted curves
and extract parameters from the fits
POLYPT accepts the output cards from LSQFIT as input, for any order polynomial,plots the experimental points( xi , yi ) and draws the curve of the polynomial, by

calculating its value at 1/20th angstrom intervals from the coefficients.
th
FLEXPT accepts the€ output cards of LSQFIT, for 4 order only, and calculates the mid-point of the line joining the inflexion points of the curve, giving an estimate of the mean
wavelength shift.
MINIPT and MINIPS accept cards from LSQFIT 4th and 6th order polynomial
respectively. They find the minimum point of the curve by successive approximation,
giving a different estimate of the wavelength shift.
GAUSPT accepts the output of GAUSSN and draws the fitted Gaussian shape.
All exit on a blank card.

221

Computer programs

222

Computer programs

223

Computer programs

224

Computer programs

Appendix 4 STREAM and CONVL
Program STREAM

For details see section 4.11.1.

Calculating Doppler shifts from a
uniform dust flow model

225

Computer programs

226

Computer programs

Program CONVL

Simulation of signal output of scanning
Fabry-Perot spectrometer

The program reads into an array A(2001-4001) a spectrum at relatively high resolution,
at 1/50th angstrom intervals. It then places in an array B(2001-4001) values of theinterference filter transmission over the corresponding wavelength range. The twoarrays are now multiplied together point by point, to place in the array E(2001-4001) the
spectrum as modified by passage through the interference filter. Using the Airyfunction with values of finesse and etalon gap read from the data, a Fabry-Perot profile
centred on the point E(3001) is put into the locations C(1-6001). The program thenconvolves arrays E and C, the sums of products for E(2001-4001) being used for shiftsof the array C from –200 to +200, the sum for each shift giving one value in the array
SUM(1-4001).

Finally the result is divided by the filter shape as widened byconvolution with the instrumental function, imitating the method used in the reductionof real scans.
At each stage a crude graphical output is printed; the process can be repeated for
different etalon parameters.
DATA:i) Spectrum – format as desired matched to a suitable read statement.
ii) Format (F2.0) FINESSE.
iii) Format (F3.0) ETALON GAP.
Blank card to exit.

227

Computer programs

228

Appendix 5

Hicks, May, Reay: NATURE paper 1972

Hicks, May and Reay NATURE paper
1972

229

Hicks, May, Reay: NATURE paper 1972

230

Hicks, May, Reay: NATURE paper 1972

231

Hicks, May, Reay: NATURE paper 1972

232

Hicks, May, Reay: NATURE paper 1972

End of Hicks, May, Reay paper, Nature 1972

233

Hicks, May, Reay: MNRAS paper 1974

Appendix 6 Hicks, May and Reay MNRAS paper
1974

234

Hicks, May, Reay: MNRAS paper 1974

235

Hicks, May, Reay: MNRAS paper 1974

236

Hicks, May, Reay: MNRAS paper 1974

237

Hicks, May, Reay: MNRAS paper 1974

238

Hicks, May, Reay: MNRAS paper 1974

239

Hicks, May, Reay: MNRAS paper 1974

240

Hicks, May, Reay: MNRAS paper 1974

241

Hicks, May, Reay: MNRAS paper 1974

242

Hicks, May, Reay: MNRAS paper 1974

243

Hicks, May, Reay: MNRAS paper 1974

End of Hicks, May, Reay paper, MNRAS 1974

244

Derivation of the Airy Summation

Appendix 7 Derivation of the Airy Summation for a
Fabry-Perot Interferometer

The diagram shows two rays emerging from the Fabry-Perot, one of which, T0 , passes
through the etalon without reflections, and the other, T1, is reflected twice by thecoatings, before emerging. The refractive index of the medium between the plates is n ,€and that of the medium outside the gap is n0 . In our case the medium outside the gap€

is the material of the plates themselves, quartz, and the emerging rays have to pass€through another refraction before they finally leave the etalon, but the pattern of the€ by this, so it is ignored in this analysis.
resulting transmission is unaffected
At each reflection, the amplitude is reduced by

R , and the phase is shifted by π ,while at each transmission through an interface the amplitude is reduced byAssuming no absorption, we have, by the law of conservation of energy, T + R =1.

€

€

€
€

T.

245

Derivation of the Airy Summation

Using complex form for a light ray, and defining the amplitude of the wave at point (a)as unity, the amplitude at point (b) is

T0 = T exp(ikl cosθ )where k is the wave number inside the etalon, given by k = 2π n λ , and λ is the€vacuum wavelength. At point (c) the amplitude is€T1 = TR exp[(2πi + 3ikl€) cosθ ]€The total amplitude of the two beams will be the sum of their amplitudes, measured at a
wavefront €defined by a line perpendicular to the direction of the beam. We therefore
need to add, to the amplitude at point (b), a modified amplitude, equal in magnitude to

T1, but retarded in phase by an amount k0l0 , where k0 is the wave number outside ofthe etalon, given by k0 = 2πn0 λ . Thus€€€
T1 = RT exp[(2πi + 3ikl ) (cosθ − ik0lo )]€where l0 is given by l0 = 2l tan θ tan θ 0€
Neglecting the 2π phase change due to the two reflections, we have for the phase€ difference between€
the two beams€

δ = 2kl cosθ − k0l0

From the law of refraction€nsinθ = no sin θ 0

€

246

Derivation of the Airy Summation

So the phase difference may be written

Δ = 2kl cosθTo within a constant multiplicative phase factor, the amplitude of the m€themergent

beam can be written as

Tm = TR m exp(imθ )€The total transmitted beam is the sum of all individual beams

€

AT =

∞
∑

m=0

= T

€

Tm

∞
∑

m=0

R m exp(imΔ)€

This is the sum of a geometric series, and so can be written as€

AT = T [1− R exp(iΔ)]∗
The intensity of the beam is now AT AT and, since the incident beam was assumed to€have an intensity of unity, this will also give the transmission function:

€
τ =

T2
1+ R 2 − 2R cos Δ

This completes the derivation of equation 2.1, page 81.

€

247

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