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The Diffuse Interstellar Medium
The diffuse interstellar medium (ISM) comprises a heterogeneous mixture of different phases of gas and dust, containing typically 10 % of the matter of the Milky Way. Observations of the ISM show that different regions can be distinguished with respect to temperature, density and ionization state of the gas. Different cloud types can be classified due to whether hydrogen, the most abundant element in the ISM, is in its neutral atomic, ionized atomic or molecular form.
A classification of cloud types was introduced by Snow and McCall [1], based on the local fraction, fXn, of the key abundant species, X, in the cloud: H/H2, C+/C/CO, whose relative abundance is indeed dependent on the ultraviolet (UV)(1) field (see Section 1.1.3). According to Snow and McCall [1], we can distinguish between diffuse, translucent, and dense molecular clouds (see Table 1.1).
Diffuse Clouds are tenuous, mostly populated by atoms, ions, and small radicals with very low gas density and exposure to the mean interstellar UV radiation field, ionizing or dissociating most of the molecules. Their typical temperatures range from 30 to 100 K. Diffuse atomic clouds contain mainly hydrogen in neutral atomic form and ionized atoms with ionization potentials, IP+, below the IP+ of hydrogen, 13.6 eV, such as carbon, providing electrons in this region. Diffuse molecular clouds, however, are characterized by a larger, more substantial fraction of molecular hydrogen, fHn2 > 0.1.
(1)The ultraviolet light ranges of the electromagnetic spectrum referred to throughout this thesis are reported in Table B.1 in Appendix B.
The UV interstellar radiation field is slightly attenuated, AV ≈ 0.2, but strong enough to ionize atomic carbon, C+, or to photodissociate CO. Carbon is therefore primarily present as C+ or as small molecules, such as CH, CO, CN, C2, or C3.
Translucent Clouds comprise less ionized atomic carbon, C+, than diffuse clouds. Carbon is mostly abundant in its neutral form or included in molecules, but the local fraction of CO remains below 0.9. They are significantly denser than diffuse clouds, and therefore the interstellar UV radiation is more attenuated. The outer edges of dense molecular clouds might be described as translucent clouds.
Dense Molecular Clouds are the densest clouds from which stars and planets can form. Inside these clouds, the typical temperatures range from 10 to 50 K. They contain predominantly molecular hydrogen and stable molecules, such as CO, but also more complex molecules are formed and survive in regions shielded from the UV field. This is due to the fact that the atoms, molecules, and submicronic particles present in the cloud efficiently absorb the UV/Visible (Vis) light emitted by nearby stars. In addition, gas and dust particles are strongly coupled, as gas-phase elements freeze onto the surface of grains. This leads to the formation of ice mantles which induce a richer chemistry [2]. In particular, in these ice mantles on interstellar grains, the formation of complex organic molecules, which have been detected in dense clouds, may be favored [3]. The interface of these clouds with hot stars gives rise to photodissociation regions (PDRs) in which the strong interaction with UV photons governs the physical and chemical conditions (see Section 1.1.3).
Cosmic Dust Life Cycle
The most abundant element of the ISM is hydrogen, making up 70.4 % of its mass, fol-lowed by helium with 28.1 %. The remaining 1.5 % are composed of heavier elements, for instance, oxygen, carbon, and nitrogen. Typically 99 % of the interstellar matter is in gaseous form and only 1 % is attributed to tiny dust particles of submicronic size. Interstellar matter lives through continual change which is closely coupled to the formation, life, and death of stars. We can thus name this continual change the cosmic dust life cycle, which is presented in Figure 1.1.
Figure 1.1: The life cycle of the interstellar matter. Credit for the original images to ESA and NASA.
Star forming regions emerge from dense molecular clouds as a natural consequence of their low temperatures and high densities. Containing enough mass to initiate a gravitational collapse, a protostar is formed. Conservation of angular momentum during the collapse causes the formation of a protoplanetary accretion disk which rotates around the star. Within the high density of the disk, coagulation of dust might lead to the formation of planets, to a stellar system. Massive stars with masses above 8 M end with explosions called supernovae, during which stellar matter is violently ejected. Evolved low mass stars with masses below 2 M and intermediate mass stars, 2 − 8 M , however, lose their mass during their life by stellar winds. The outer layer of such asymptotic giant branch (AGB) stars form expanding circumstellar envelopes and, during the final phases of their evolution, planetary nebulae. The progressive mass loss causes matter being constantly ejected by the stars, feeding the interstellar medium and enriching the interstellar matter which will eventually form new molecular clouds.
Photodissociation Regions
Photodissociation regions (PDRs), sometimes also called photon-dominated regions, are regions of the ISM in which the physical and chemical conditions are governed by the interaction of matter with far UV (FUV)(2) photons. In these regions, hydrogen is found in its neutral form and only photons below 13.6 eV are present, the so-called H i regions. Bright PDRs often result from the interaction of young massive stars with their parental cloud [4]. Typical examples are the Orion Bar or the north-west PDR of NGC 7023, for which a recent study has gathered all available observational data to obtain a better understanding of the physical conditions in these environments [5]. These bright PDRs are usually found at the interface with H ii regions that surround massive stars and in which hydrogen is ionized and photons above the 13.6 eV limit can therefore penetrate.
The FUV flux in PDRs is usually quantified in units of G0 given by Habing [6]. It is normalized to the value corresponding to the average interstellar flux integrated between 5.2 and 13.6 eV, which results in an energy flux of 1.6 • 10−10 W cm−2 and a photon flux of 108 photons cm−2 s−1. In massive star forming regions, values of G0 of 104 − 106 can be found. All of the atomic and most (90 %) of the molecular gas in the Galaxy is contained in PDRs [4]. The interaction of FUV photons with the gas and dust impacts both the thermal balance of the gas and the chemistry, leading (2)We consider here UV photons up to 13.6 eV.
Figure 1.2: Schematic of a typical photodissociation region. The PDR is illuminated by UV photons from the left and extends from where atomic matter dominates to where the photodissociation of O2 is negligible, at AV ≈ 10. This is where, what we consider, the molecular region starts. Credit for the original image of the Orion Nebula to Hubble/ESA/NASA.
to a gas composition which depends on the penetration of the FUV field. H/H2 and C+ /C/CO interfaces are predicted by models [see 4]. A schematic of a typical PDR is shown in Figure 1.2. Typically one percent of the mass in PDRs is in the form of dust particles (see Section 1.1.2). Dust plays a key role in PDRs by attenuating the FUV field and generating electrons by the photoelectric effect which heat the gas [7, 8]. It also provides surfaces on which molecules adsorb, favoring chemistry and increasing the possibility of chemical reactions. A well-known case is the formation of the most abundant molecule, H2 , as described in a recent review by Wakelam et al. [9], but this concerns also complex organic molecules [3]. The physicochemical processes in PDRs are complex and coupled with each other. PDR models have been developed to analyse the emission in the gas lines and to constrain the chemical and physical processes that take place in PDRs [10–14].
Dust Extinction and Emission
When observing the night sky in the visible (Vis) light spectrum range as presented in the top panel of Figure 1.3, dark features are observed in the galactic plane. These dark spots are attributed to cosmic dust in the diffuse ISM, efficiently extinguishing starlight by absorbing and scattering all UV/Vis photons emitted by hot stars, leading to dust extinction [15]. Changing the observational technique to the infrared (IR), in particular to the mid (MIR) and far infrared (FIR), bright emission features are observed where it was dark before (see Figure 1.3, bottom panel). Hence, we conclude that most of the absorbed energy is emitted back to the ISM in the IR.
The left panel of Figure 1.4 shows the dust extinction curve of the diffuse ISM as given by Compiègne et al. [16]. It represents the extinction cross section expressed per H nuclei as a function of inverse wavelengths, λ−1, i.e., wavenumber, ν˜, in µm−1. Its main characteristics [17] are the gradual rise from the IR to the UV range and the bump at approximately 220 nm (4.55 µm−1 or 5.64 eV) which is attributed to π → π∗ electronic transitions in polycyclic aromatic hydrocarbons and very small graphite-like grains [18, 19]. In addition to the main dust features, a set of narrower absorption bands have been observed in several lines of sight when investigating reddened early type stars [20]. These are called diffuse interstellar bands (DIBs) and were first discovered one century ago [21, 22]. The number of observed DIBs has increased over the years, counting over 400 detected DIBs [23, 24]. Arising at optical and near infrared wavelengths in the interstellar gas, the DIBs are thought to be associated with carbonaceous molecules, but their origin has not yet been conclusively revealed. Recently, a couple of bands in the near IR could be attributed to originate from the cationic buckminsterfullerene, C+60 [25, 26]. Being an ongoing unsolved mystery in astronomical spectroscopy, new DIBs are discovered regularly. See Geballe [27] for a detailed review on the recent developments on DIBs. Dust reemits the energy it has absorbed in the Vis/UV range in the IR to mm range. The characteristics of the emission curve depend on the composition of the interstellar dust population, i.e., the chemical nature of the dust and its size. The spectral energy distribution of dust emission in the galactic interstellar medium has been modeled by Compiègne et al. [16], attributing the various emission features to different kinds of dust as depicted in the right panel of Figure 1.4. Different dust models have been built to account for the features of the extinction and emission curves [16, 28–31]. These concentrate mainly on dust made of silicates and carbonaceous grains of varying sizes, typically in the molecular to µm size. Characteristic features in the low wavelength range can be seen from the right panel in Figure 1.4, especially at 3.3, 6.2, 7.7, 8.6, and 11.3 µm. Initially, they were referred to as the Unidentified Infrared Emission (UIR or UIE) bands and were attributed to very small grains (VSGs) with radii in the ∼10 − 150 Å range [32]. After the comparison to laboratory data, another potential carrier was proposed [33, 34].
Astrophysical Polycyclic Aromatic Hydrocarbons
Figure 1.5: The Spitzer-IRS spectrum of the nuclear region of the galaxy NGC 4536 and the ISO-SWS spectra of the planetary nebula NGC 7027 and the photodissociation region at the Orion Bar illustrate the richness and variety of the PAH spectrum. Also indicated are the vibrational mode identifications of the major PAH bands. Taken from Peeters [35].
Table of contents :
1 Scientific Context and Objectives
1.1 Astrophysical Background
1.1.1 The Diffuse Interstellar Medium
1.1.2 Cosmic Dust Life Cycle
1.1.3 Photodissociation Regions
1.1.4 Dust Extinction and Emission
1.1.5 Astrophysical Polycyclic Aromatic Hydrocarbons
1.1.5.1 Aromatic Infrared Bands
1.1.5.2 Characteristics of Astro-PAHs
1.2 Photophysics of Astro-PAHs
1.3 Laboratory Astrophysics Studies on PAHs
1.3.1 Spectroscopy of PAHs
1.3.2 Studies on the Photophysics and Chemistry of PAHs
1.4 Objective of this Work
1.5 Implementation within the EUROPAH Network
2 Laboratory Methods
2.1 Ion Traps and Storage Devices
2.1.1 Production of Ions
2.1.1.1 Atmospheric Pressure Ionization Source
2.1.1.2 Electron Impact Ionization
2.1.1.3 Laser Desorption Ionization
2.1.2 Trapping and Storage of Ions
2.1.2.1 Paul Trap
2.1.2.2 Penning Trap
2.1.2.3 Electrostatic Storage Rings
2.2 Mass Spectrometry Techniques
2.2.1 Mass Selection and Resolution
2.2.2 Quadrupole Mass Spectrometry
2.2.3 Fourier Transform Ion Cyclotron Resonance Mass Spectrometry
2.3 Spectroscopy Techniques in Ion Traps
2.3.1 Action Spectroscopy on the Bare Cation
2.3.1.1 Single Photon Absorption
2.3.1.2 Multiple Photon Dissociation
2.3.1.3 Two Photon Absorption
2.3.2 Predissociation Spectroscopy of the Tagged Ion
2.3.3 Light Sources
2.3.3.1 Tabletop Light Sources
2.3.3.2 Synchrotron SOLEIL
2.3.3.3 Free Electron Laser FELIX
3 Theoretical Methods
3.1 Theoretically Studying Astro-PAHs
3.2 Density Functional Theory
3.2.1 Born-Oppenheimer Approximation
3.2.2 Electron Density
3.2.3 Exchange-Correlation Functional
3.2.4 Basis Sets
3.2.5 Potential Energy Surface
3.2.6 Counting States
3.3 Time-Dependent Density Functional Theory
4 Large PAH Cations under VUV Irradiation
4.1 Introduction
4.2 Beamline DESIRS with LTQ Ion Trap
4.3 Experimental Methods and Data Analysis
4.3.1 Acquisition of Mass Spectra
4.3.2 Photon Flux Calibration
4.3.3 Detector Gain Efficiency
4.4 Results and Discussion
4.4.1 Action Spectra
4.4.2 Competition between Ionization and Dissociation
4.4.3 Theoretical Photoabsorption Cross Sections
4.4.4 Experimental Photoproduct Cross Sections
4.4.4.1 Approach 1 – Introduction of a Proportionality Factor
4.4.4.2 Approach 2 – Determination via a Known Cross Section
4.4.5 Photoionization Yields
5 Isomer Diversity in the –H Fragments of Methylated PAH Cations
5.1 Introduction
5.2 Motivation
5.3 FELion Beamline at the FELIX Laboratory
5.4 Preparation of the Experiment
5.4.1 Methylated PAH Precursors
5.4.2 Expected Results: IR Spectra
5.5 Data Acquisition and Analysis
5.6 Results and Discussion
5.6.1 –H Fragment of 1-Methylpyrene: C17H11
5.6.2 –H Fragment of 2-Methylanthracene: C15H11
5.6.3 –H Fragment of 2-Methylnaphthalene: C11H9
5.6.4 Isomer Diversity for C11H9
5.7 Evolution with Size
5.7.1 Summary of Results
5.7.2 Toward Changing the Isomer Mixture
6 Astrophysical Relevance and Perspectives
6.1 Astrophysical Implications
6.1.1 Charge State of Astro-PAHs
6.1.2 Photoionization Yield for Astrophysical Modeling
6.1.3 Methyl and Methylene Sidegroups attached to PAHs
6.2 Laboratory Astrophysics Perspectives
