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Early-type stars, cosmic rays and magnetic fields
Embedded stars were clearly a source of energy: Str¨omgren (1939) showed that early-type stars will form ionization fronts, expanding because of overpressure and driving a shock wave into the interstellar medium but only a tiny fraction of this energy is converted into kinetic energy, so other sources had to be searched.
Cosmic rays are energetic charged particles, mostly protons, originating from outer space. The term ray is historical as cosmic rays were thought to be electromagnetic radiation. Fermi (1949) suggested that the cosmic rays were accelerated primarily in the interstellar space of the galaxy by collisions against moving magnetic fields. Relative dense clouds with a root-mean square motion of about 30 km/sec create streaming motions in the tenuous and ionized intercloud medium (0.1 atom/cc). Magneto-elastic waves will convert the kinetic energy into magnetic energy (Alfven, 1942) and build up a magnetic field (about 5 µG, or 0.5 nT). The lines of force will form a crooked pattern since they will be dragged in all directions by the interstellar matter.
Fast particles (a few GeV, so close to the velocity of light) will spiral around these lines of forces until it collides against an irregularity in the cosmic field (region of high intensity field or curved line of force) and so are reflected. If the magnetic field is slowly variable, the particle will gain or lose energy after the reflection. The net result will be an average gain, because head-on collisions (energy gain) are more frequent than overtaking collisions (energy loss). This relatively simple model yields an inverse power law for the spectral distribution of the cosmic rays. A comparison with data then available gives a mean distance between collisions of the order of a light-year.
The existence of a magnetic field in the interstellar medium was confirmed by po-larization observations in the direction of distant arms. The polarization is due to a magnetic orientation of the interstellar dust particles, leading to different amounts of absorptions of light polarized parallel and perpendicular to the magnetic field. The first maps of the polarization effect showed that the galactic magnetic field was roughly parallel to the direction of the local arm.
By considering the velocity of magneto-hydrodynamic waves, Chandrasekhar & Fermi (1953) demonstrated that the magnetic field is inversely proportional to the angular deviation between the plane of polarization and the direction of the spiral arm, deducing an estimate of 7.2µG for the field intensity. With an independent method, they derived an estimate of 6µG, using the requirement of equilibrium of the spiral arm with respect to lateral expansion and contraction.
History of the discovery
The solar wind has a direct and visible impact on Earth: aurorae have been reported in the ancient literature from both East and West. In 1731, the French philosopher de Mairan proposed that the aurora was connected to the solar atmosphere, he suspected a connection between the return of sunspots and the aurora. In 1859, Carrington and Hodgson observed independently a solar flare that was followed by a geomagnetic storm the day after. Carrington suspected a connection between both events and suggested the existence of a continuous stream of particles flowing outward from the Sun.
Around 1916, Birkeland showed with geomagnetic surveys that auroral activity was nearly uninterrupted, he concluded that the Earth was continually bombarded by charged particles emitted by the Sun. Chapman proposed that the geomagnetic storm is the result of a coherent cloud ejected from the sun with a thousand km/s velocity at the time of a solar flare. In the early 1950s, Biermann pointed out that the observed motions of comet tails would seem to require gas streaming outward from the Sun. He suggested that gas is often flowing radially outward in all directions from the Sun with velocities ranging from 500 to 1500 km/s.
Eugene N. Parker showed that the solar wind originates from the hot corona (about 2 MK) that expands radially into interplanetary space (Parker, 1958). This flow becomes supersonic at a few solar radii. The spacecraft Luna 2 and Mariner II brought experimental confirmation with measurements of the plasma parameters (Neugebauer & Snyder, 1962).
Distortion by the interstellar magnetic field
Fahr et al. (1986, 1988) developed the first MHD model of the heliopause (HP) taking into account the interstellar magnetic field for any orientation. Pressures on both sides of the HP can be described by the Newtonian approximation (NA). For this purpose, the unperturbed full MHD stress tensors of the two plasma flows are projected onto the normal to the surface, and yield a quadratic partial differential equation in the general 3-dimensional case. The effects of the interplanetary magnetic field are neglected since the magnetic pressure is negligible. If solar wind asymmetries are neglected too, the plane containing the LISM velocity vector and the interstellar magnetic field vector ({ V IS , B IS }) becomes a main symmetry plane of the heliopause. Deriving the point of the maxi-mal total pressure, Fahr et al. (1988) found a symmetry axis contained in the plane V , B {→−IS →−IS } and deviating from the upwind direction by an angle: 2 MA2 − cos(2ψ0) θ0 = 1 atan −sin(2ψ0) where MA is the Alfven number in the interstellar medium and ψ0 the angle between the LISM velocity vector and the interstellar magnetic field vector. The presence of a magnetic field in the interstellar medium induces a tilt in the orientation of the heliopause.
Velocity variations of the interplanetary hydrogen
During the last two decades, SOHO/SWAN and the echelle modes of the Goddard High Resolution Spectrograph (GHRS) and Space Telescope Imaging Spectrograph (STIS) on the Hubble Space Telescope (HST) have been used to measure the Ly-α Doppler shift with respect to the heliospheric referential and line profile with greater precision than previous observations (Clarke et al., 1998; Scherer et al., 1999; Ben-Jaffel et al., 2000; Qu´emerais et al., 2006). As shown in Table 2.1, IPH spectroscopic observations cover several decades, and while not uniformly spaced in time, they roughly span the entirety of solar cycle 23.
Table of contents :
Contents
R´esum´e
Introduction
1 Interaction between the solar wind and the interstellar medium
1.1 Exploring the interstellar medium
1.1.1 Perspectives: the Milky Way and beyond
1.1.2 Detection of the interstellar matter
1.1.3 Early-type stars, cosmic rays and magnetic fields
1.1.4 Supernovae and stellar winds
1.1.5 The chimney model
1.1.6 The Local Cavity
1.1.7 The Local Interstellar Cloud
1.2 The solar wind
1.2.1 History of the discovery
1.2.2 The Sun
1.2.3 Models of solar wind
1.2.4 The heliosphere
1.3 Plasma models and structure of the heliospheric interface
1.3.1 Orders of magnitude
1.3.2 The two-shock model
1.3.3 Distortion by the interstellar magnetic field
1.4 The interplanetary hydrogen
1.4.1 Cold and hot models
1.4.2 Charge exchange in the heliospheric interface
1.4.3 Effects of solar activity
1.4.4 Influence of the interstellar magnetic field
1.4.5 Radiative transfer of Lyman-α photons
1.5 Summary
2 Analysis of the bulk velocity of the interplanetary hydrogen
2.1 Introduction
2.1.1 The heliospheric interface in a nutshell
2.1.2 Modification of the interstellar hydrogen flow
2.1.3 Velocity variations of the interplanetary hydrogen
2.2 Observations with HST
2.2.1 Instruments: GHRS and STIS
2.2.2 Description of the signal
2.2.3 Contamination in STIS observations
2.3 Data Analysis
2.3.1 Fitting procedure
2.3.2 Error analysis
2.3.3 Results
2.4 Discussion
2.4.1 Comparison with other studies
2.4.2 Possible reasons for discrepancies
2.4.3 The need for new data
2.5 Conclusion
3 HYPE-INSPIRE
3.1 Scientific motivations
3.1.1 Scientific goals
3.1.2 Benefits of interferometers in astrophysics
3.2 All-reflective Spatial Heterodyne Spectrometer (SHS)
3.2.1 Genesis and development
3.2.2 Instrument design
3.2.3 Interference fringes
3.2.4 Resolving power, field of view and bandpass
3.2.5 Anti-aliasing
3.2.6 Intensity distribution
3.3 Polarimetry measurements
3.3.1 State of the art
3.3.2 Basics of polarimetry
3.3.3 Design of the polarimeter
3.4 Payload design
3.4.1 Optical layout and imaging
3.4.2 Detector
3.4.3 Mechanical structure and electronics
4 Experimental challenges
4.1 The vacuum-ultraviolet range
4.1.1 Choice of optics
4.1.2 Working under vacuum
4.1.3 Constraints in space applications
4.2 Efficiency anomaly on a diffraction grating
4.2.1 Context
4.2.2 Generalities on gratings
4.2.3 Grating anomalies and VUV region
4.2.4 Experimental procedure
4.2.5 Numerical simulations
4.2.6 Discussion and conclusion
4.3 Next steps for HYPE-INSPIRE
4.3.1 Integration
4.3.2 Flight window
Conclusions
A Plasma physics
A.1 Electromagnetism: Maxwell’s equations
A.2 Kinetic theory of gases: the Boltzmann Equation
A.2.1 Statistical description of particles
A.2.2 Macroscopic quantities
A.2.3 The Boltzmann Equation
A.3 The fluid equations
A.3.1 Preliminary
A.3.2 The particular case of elastic collisions: the ideal gas
A.3.3 General case
A.4 (Magneto)hydrodynamics
A.4.1 The speed of sound
A.4.2 Shock physics: the Rankine-Hugoniot relations
A.4.3 Frozen magnetic field lines
A.5 The kappa distribution
B Ultraviolet spectrometry with HST
B.1 Description of the instruments
B.1.1 GHRS
B.1.2 STIS
B.2 Sensitivity at Ly-α
B.2.1 Simulation of the throughput
B.2.2 Calculation of the sensitivities
B.3 Exposure times of the IPH observations
C Fundamental constants
D Acronyms
Bibliography
