USOS DEL JENGIBRE

Sgr A* is detectable in the radio, NIR and X-ray regimes of the electromagnetic spectrum; the spectral energy distribution is shown in the left panel of Fig.1.15. The compact radio source is detectable at all times; its spectrum is roughly described by a powerlawνLν ∝ν1.3,

which rises up to submm wavelengths (submm peak has a luminosity of about 5×1035 erg s−1_{), whereafter the spectrum suddenly drops down to less than the detection limit in}

the NIR regime. The other not ambiguous detection of steady emission is in the X-rays, where the steady faint X-ray emission has a luminosity of∼1033erg s−1 in the typical X-ray band 2-10 keV (Baganoff et al., 2003). In the left panel of Fig.1.15 is also shown a model for the quiescent emission; the radio spectrum is well described by optically thick synchrotron emission from relativistic thermal electrons, with a temperature of a few 1010K (Lorentz

γ factor ∼10) and a density of 106_cm−3_{, spiralling around a magnetic field of about 20-50}

Gauss.

The quiescent X-ray emission is thought to arise from thermal bremsstrahlung origi- nating from the transition region between the accretion flow and the surrounding ambient medium. The second maximum in the model of the quiescent emission, at a frequency around 1016 Hz, is due to inverse Compton scattering of the synchrotron photons by the same population of thermal electrons. As we can see from the SED of Sgr A*, the luminos- ity of the source over the entire electromagnetic spectrum is very low when compared to the luminosities of SMBHs in other Galaxies. In the X-rays, for example, the measured qui- escent state luminosity is 1033_{erg s}−1_{, about 10}−11 _{the Eddington luminosity for a SMBH}

with the mass of Sgr A*. This low luminosity is not due to the lack of gas to accrete; Cuadra et al. (2006) stated that there is enough gas from stellar winds in the surroundings of Sgr A* which could produce a much enhanced radiative output, up to 105 _{times the}

measured one. The very low activity is instead due to a radiatively inefficient accretion flow onto the SMBH (for theoretical models see, e.g., Narayan et al., 2002; Quataert, 2003). The steady emission from Sgr A* shows different polarisation properties within differ- ent energy bands. At frequencies higher than 100 GHz (submm) the emission is linearly polarised at a level of about 10%, whereas in the radio regime the linear polarisation is very low and the emission has a circular polarisation on the level of 0.3-1% (Bower et al., 1999). Due to the lack of X-ray polarimeters in space, there is no observation yet of the polarisation of X-ray emission from Sgr A*, as well as from other X-ray sources.

Sgr A* also shows highly variable emission. Many observational campaigns over the last decades in the radio, submm, NIR and X-ray bands have discovered that this peculiar source is variable over the entire observed spectrum, with the degree of variation vary-

Figure 1.15: Left: spectral energy distribution of the steady state of Sgr A* across the electromagnetic spectrum (Genzel et al., 2010). All numbers are given for a galactocentric distance of 8.3 kpc (Gillessen et al., 2009). Right: X-ray (top), NIR and submm (middle), and radio (bottom) lightcurve of Sgr A* over the April 4, 2007 flare Dodds-Eden et al. (2010).

ing dramatically between different regimes. At radio wavelengths, the ratio between the maximum and minimum detected fluxes is of the order of one (cm wavelengths) to a few (submm); on the other hand, in the NIR ad X-ray bands, the flux excursion between the quiescent minimum level and the brightest peak is significantly higher and goes from a factor of 20 in the NIR band to a factor of ∼160 in the X-rays, where the brightest flare has been measured to have an X-ray luminosity of 3.9×1035 erg s−1 (Porquet et al., 2003). Flares in the NIR occur on average about 4 times per day (Eckart et al., 2006), with a typical duration of ∼80 min (Genzel et al., 2010). Contrary to the NIR flares, X-ray ones are detected less frequently, with an average of about 1 per day (Baganoff et al., 2003), with typical duration around 50 min.

In the last decade several models for the flare emission have been developed. The spatial coincidence within few milliarcsec of the NIR flares and the central massive BH and the short timescales of the variability suggest that the flares are generated in the inner region of the accretion flow; this excludes tidal disruption of a star, for example, which would otherwise make the flare much longer than observed. An increase of the accretion rate of the radio-submm source is also unlikely, since the excursion observed in the X- ray luminosity would require a similar excursion in the radio-submm range, which is not seen (Markoff et al., 2001). Baganoff et al. (2001) and Markoff et al. (2001) suggested that the flare might be generated in response to a fast acceleration of electron within the innermost regions of the accretion flow, via magnetic reconnection events and/or accretion instabilities). The accelerated electrons, which can have γ factors up to 103_{, could then}

1.3 The super massive black hole Sgr A* 39

upscatter NIR or submm photons to the X-rays via the Inverse Compton effect; besides this, X-ray emission in flares can also be produced by synchrotron emission of highly energetic (γ ∼106_{) electrons in a magnetic field of around 5 Gauss (Dodds-Eden et al., 2010).}

Whatever the emission mechanism for the NIR and X-ray flares is, the luminosity reached by Sgr A* during these sudden episodes of brightening is still very low compared to normal active galactic nuclei in other galaxies. The question that consequently arises and is a subject of ongoing research is whether Sgr A* might have been much brighter in the past, and is currently in a resting low activity phase. A recent intriguing discovery is that of the so called Fermi bubbles (e.g. Su et al., 2010), two extreme structures about 15 kpc parsec in size with the shape of two giant bubbles extending north-south, symmetrically with respect to the galactic plane; these structures, so far only detected in the γ energy band, might be the remnant of a past catastrophic event related to Sgr A*, although the discovery is very recent and the research on the nature of these gigantic diffuse structures has just begun. On the other hand, for what concerns closely the work presented in this thesis, there is also the hypothesis that Sgr A* could have undergone a period of higher activity, with an X-ray luminosity estimated around 1039_{erg s}−1 ₍∼₁₀−5_L

Edd) having occurred some

hundreds years ago. We will discuss the origin of this idea and the observations carried out to test it in Chapter 3.

Chapter 2

The physics behind

In this thesis work, we have been mainly dealing with Fe-Kαline emission from neutral Fe in the massive MCs in the GC region. In this chapter, we will briefly review the chemical and physical properties of the element Fe (iron), and then describe the ionisation pro- cesses produced by the absorption of (UV-X-rays) photons and by particle bombardment (electrons and ions). To conclude, we are going to describe the phenomenology of the two physical processes involved in the Fe-Kα line emission, discussing the X-ray spectra of an MC either illuminated by X-ray photons or bombarded by low energy CR electrons and protons. However, we must warn the reader that some of the physical processes described in this Chapter are not directly observable (and therefore are not part of our observa- tional doctoral work) with current X-ray observatories because of the insufficient spectral resolution of current instruments; nevertheless, a full theoretical understanding of the Fe ionisation process gives a unique background for future observations of this line emission, whose hyperfine effects (in part described in this Chapter) will be measured by the next generation X-ray detectors (i.e. X-ray calorimeters) in less than a decade from now.

2.1

Why Iron?

Why is Fe special? In the 1-10 keV energy range, the one observed by current operating X-ray satellites like XMM-Newton, Chandra and Suzaku, the Fe-Kα lines from different ionisation stages of Fe are the strongest features one can see, and this is because Fe is the element with the highest (by a factor of ∼5) product of abundance and fluorescence yield (for the K-shell). The most abundant stable isotope of Fe is 56_{Fe, whose atomic nucleus}

is composed of 26 protons (atomic number Z=26) and 30 protons, with an atomic mass measured to be 55.845 atomic mass unit 1.

In Fig.2.1 we show the electronic configuration of the Fe atom, which can also be written as [Ar]3d6_4S2_{, where [Ar] is the electronic configuration of Argon, i.e. 1S}2_2S2_2P6_3S2_3P6_{. In}

the right panel of this figure we show the energy levels of the inner shell electrons and the electronic transitions allowed to fill a K-shell vacancy. In this plot, n is the fundamental

Figure 2.1: Atomic structure (Robson, 2011) and energy levels (Phiwe , 2008) of the K-L-M shells of Iron.

quantum number of the atomic structure, which identifies the zero order energy of an electron bound to the nucleus (and therefore a negative energy). The second quantum number, l, refers to the angular momentum of the electron in a certain shell; as a rule,

l can assume n values between 0 and n-1. The last quantum number is j, which is the quantum number which describes the total angular momentum of the electron once the spin orbit coupling has been taken into account; the permitted values ofj are l±1/2. Whenever an inner shell vacancy occurs, the atom is in an excited state, and therefore tends to ”relax” by filling this vacancy with an outer shell electron; subsequent to this transition an emission of either a photon or an outer shell electron occurs (see later), because the electron energy is now lower than in the original state (inner shell electrons are more bound to the nucleus). The quantum mechanical selection rules state that the only allowed transitions are those for which ∆l=±1 and ∆j=0,±1 (see, for example Eisberg & Resnick, 1985). Therefore, the Fe-Kα line from neutral Fe can occur in two separate ways, defined Kα1 and Kα2, in

which the upper level from where the electron originates is described by l=1 and j=1/2 andj=3/2, respectively. Because of the rule ∆l=±1, the transition L1 →K is not observed

(see Fig.2.1). The energies of the Fe doublet are 6.391 keV and 6.404 keV, respectively for the Kα1 and Kα2 (X-ray data booklet , 2009).