Large-scale emission

Although jets and lobes are characteristic of powerful radio galaxies, these structures, on a smaller scale, can also be found in objects that are not strictly considered radio- loud (according to the definition in eq. 1.1). Examples of these structures in elliptical

Figure 1.5: Example of radio emission in a low power source. Image fromCroston et al.(2007), illustrating the match between X-ray and radio emission in NGC 3801. The lobes in this source are believed to be driving a shock into the ISM.

galaxies can be found in e.g. NGC 3801 (Croston et al. 2007, and Fig.1.5), NGC 1062 (Kadler et al. 2004) and Markarian 6 (Chapter3). The effect of jets and lobes on the host galaxy is even more interesting in the case of Seyferts, which are typically hosted by spirals (see e.g. Croston et al. 2008b;Hota & Saikia 2006;Gallimore et al. 2006, and Chapter4), where the effects on dynamics and star formation are potentially much larger (see Section1.2).

The classification ofFanaroff & Riley(1974) divides radio galaxies into two sub- groups. FR I objects are core-brightened, and exhibit lobes, while FR II objects are edge-brightened. Although this is a purely morphological classification, Fanaroff & Rileyshowed that the luminosity of the source is related to its morphology, so that FR I galaxies (typically) have L178MHZ ≤ 2 × 1025 W Hz−1 sr−1. Ledlow & Owen (1996)

showed that the division between both classes depends on the optical luminosity den- sity of the host galaxy, so that the break rises with luminosity as L1.7_opt. The boundary can also be defined in terms of the long-term average jet power, ¯Q ∼ 5 × 1043 erg s−1

(Rawlings & Saunders 1991).

Both classes of radio galaxy have common structures: jets (which may or may not be visible, the latter case being more common for FR II sources) and lobes. Some

1.1. ACTIVE GALAXY STRUCTURES AND THEIR PROPERTIES 19

examples of the different morphologies can be seen in Fig. 1.6. It is important to note that the FR I / FR II division does not correspond to the high/low-excitation one (Hardcastle et al. 2009). While most low-excitation objects seem to be FRI, there is a population of bona-fide FRII LERGs, as well as numerous examples of FRI HERGs (e.gLaing et al. 1994). This is most likely caused by the complex underlying relation between fuelling, jet generation and environmental interaction.

Figure 1.6: Mosaic displaying radio maps (Morganti et al. 1993, 1999) for a variety of sources in the 2Jy sample. Top left: PKS 0034-01 (3C 15), an FR I LERG with a classic radio morphology and a jet. Top right: PKS 0349-27, an FR II NLRG. Bottom

left: PKS 1648+05 (Hercules A) an FR I LERG in a very dense cluster environment,

that is collimating the jet near the central AGN. Bottom right: PKS 2211-17 (3C 444), a somewhat atypical FR II LERG in a dense cluster environment, into which the lobes are driving a shock (Croston et al. 2011).

Jets and hotspots

Jets are defined as collimated, relativistic outflows that are produced near the central black hole and extend to scales ranging from kpc to Mpc. The base of the jet is thought to be very similar for FR I and FR II galaxies, and highly relativistic, which sometimes produces an apparent superluminal motion (see e.g. Pearson 1996;Zensus 1997). The jets of FR I are generally shorter, and decelerate to sub-relativistic speeds, presumably due to entrainment of the environment (see e.g. Bicknell 1984;Laing & Bridle 2002). The large opening angle, and the presence of knots (e.g. Hardcastle et al. 2007b;

Jorstad et al. 2005), suggest that the energy transport to large scales is not very efficient in these objects. By contrast, the jets of FR II remain collimated and relativistic for much larger distances (typically of the order of 102kpc, see e.g. Hardcastle 2009), and are more efficient at transporting the energy. FR II jets typically terminate in hotspots, bright compact regions where particle acceleration occurs, which results in a flat radio spectral index (0.5 < α < 0.7).

The composition of the jets is not well understood. We know that they contain electrons, possibly positrons (e.g. Reynolds et al. 1996; Dunn et al. 2006), and that magnetic fields are involved (see e.g. Pudritz et al. 2012, for a review). Observations seem to indicate that B fields are initially parallel to the jet axis, and that interaction with the environment and shocks alters this, introducing a perpendicular component (Attridge et al. 1999). This may be the reason why perpendicular fields are more often detected in FR I galaxies in polarization measurements (Reichstein & Gabuzda 2012). The emission mechanism that allows us to see these jets is synchrotron radiation, which is emitted when a charged particle travelling at relativistic speeds interacts with a magnetic field (in the non-relativistic regime the emission is called cyclotron). The particle moves in a helicoidal trajectory, with a characteristic orbital frequency given by (Rybicki & Lightman 1986):

ωB=

qB

γmc (1.21)

where q is the charge and γ is the Lorentz factor (related to the total energy of the particle by E = γmc2). For an isotropic distribution of velocities, the particle loses energy at a rate given by eq. 1.23. The spectrum of synchrotron radiation spans a broad range of frequencies for a given particle energy, but it is narrowly peaked around

1.1. ACTIVE GALAXY STRUCTURES AND THEIR PROPERTIES 21

a critical frequency, given by (in the case of an electron, see e.g. Longair 1992):

νc =γ2

eB

2πme

(1.22)

Since the energy loss time scale is proportional to γ−1 _{and E}−1_{, an X-ray emitting}

electron loses energy ∼ 105_{times faster than a radio-emitting electron, and protons of}

a given energy lose their energy ∼ 106times more slowly than electrons. The synchrotron power emitted by each electron is given by:

PS = 4 3σTcγ

2_β2_U

B (1.23)

where UB is the magnetic energy density, UB = B2/2µ0 (µ0 is the vacuum perme-

ability). This has the consequence that the ratio between PC (the inverse-Comptom power) and PS is the same as the ratio between that of the magnetic and photon energy densities. Another consequence of this is that the cooling of the electrons via inverse- Compton processes is very efficient, which limits the brightness temperature of radio sources to ∼ 1012_{K (this effect is often called the Compton catastrophe).}

In radio galaxies the electron population emitting synchrotron radiation has a broad range of Lorentz factors. To obtain the power law spectrum that we see at radio fre- quencies, the electrons must also have a power law distribution in energy, given by

N(E) = N0E−pover a given energy range. The total emissivity for the electron popu-

lation is then given by:

J(ν) = Z Emax Emin j(ν)N(E)dE = Z Emax Emin j(ν)N0E −p_{dE (1.24)}

where j(ν) is the emissivity of a single electron, which is a function of the critical frequency.

If we assume that j(ν) is a narrow function, and deriving E from eq. 1.22, the total emissivity is: J(ν) ∝ N0ν −p−1 2 B p+1 2 (1.25)

The shape of the synchrotron spectrum is often described as a function of the radio spectral index α, given in terms of the flux density as S (ν) = S0ν−α. From above,

The relativistic nature of the jet emission causes a boost in the luminosity of the core in radio-loud sources with low orientation angles (BL-Lacs, BLRGs and QSOs), due to relativistic Doppler beaming. This is studied in Chapter6for a range of sources. Given the rapid energy losses at higher frequencies, synchrotron emission from jets and hotspots can only be observed in X-rays if there is a source of particle acceleration (Hardcastle et al. 2007c). See e.g. Worrall(2009) for a review of X-ray jets.

Lobes

Although some of the dynamics and energetics of FR I and FR II radio lobes are very different, many assumptions can be applied to both categories.

Lobes are created when the jet interacts with the surrounding environment (ISM, IGM at larger scales), creating bubbles that expand through this medium and are filled with the particles that travel up the jet. The early models ofScheuer(1974) proposed a variety of scenarios for the propagation of the lobes through the interstellar medium of the host galaxy. If the lobes are always overpressured with respect to their surround- ings, they expand supersonically, driving a bow shock, and are not influenced by the external medium (other than for their expansion speed). Most of the analytical models describing the dynamics of radio lobes are based on this assumption (e.g. Begelman & Cioffi 1989;Kaiser & Alexander 1997), and this is the case for the sources studied in Chapters3and4. However, this is not always so, in many cases, and particularly as the lobes evolve, the pressure inside and outside the lobes becomes roughly equal, and the situation becomes more complicated (e.g. Alexander 2002;Hardcastle & Krause 2013).

The particle content of FR I and FR II lobes is thought to be quite different: the former are dominated by non-radiating particles (Croston et al. 2008a), while the oppo- site applies for FR II (Croston et al. 2005). This has implications for the overall energy budget of the lobes, and means that, a priori, the same correlations between jet kinetic energy and radio luminosity cannot be applied across both populations (although the dependence of radio luminosity with environment compensates for this fact to some extent, seeHardcastle & Krause 2013;Godfrey & Shabala 2013, and Chapter6).

As such, the emission inside the lobes of FR I sources is synchrotron in the radio regime, but thermal emission from the shocked shells dominates at X-ray frequencies,

1.1. ACTIVE GALAXY STRUCTURES AND THEIR PROPERTIES 23

while in FR II the higher electron content (and larger scale) makes inverse-Compton the dominant mechanism for X-ray emission.

The minimum energy stored in the lobes can be calculated if the magnetic field strength and particle distribution are known, asBurbidge(1956) showed that the mini- mum energy condition is very close to equipartition. If we assume equal filling factors in the radiating particles and magnetic fields, equipartition is given by:

UB = (1 + κ)

Z Emax

Emin

EN(E)dE (1.26)

where κ is the proton to electron energy ratio, and UB is the magnetic energy density,

UB = B2/2µ0.

The main issue with this assumption, however, is that the magnetic field cannot be directly determined from synchrotron emission alone. Inverse-Compton emission in X-rays can be used to constrain the electron density, and used in conjunction with the radio observations to determine the magnetic field strength (see e.g. Croston et al. 2005). X-ray emission can also be used to test how far from equipartition the lobes are. We now know that in FR I sources the thermal contribution to the energy budget cannot be neglected, since the equipartition pressure is much lower than the pressure of the gas outside the lobes (Morganti et al. 1988). In FR II systems deviations from equipartition are not so drastic (Croston et al. 2005), though they are difficult to measure in some cases, due to the limitations of geometrical assumptions for sources with low projection angles.

The main sources of photons for inverse-Compton scattering in radio lobes are the cosmic microwave background, the AGN and stellar light. Of these, AGN photons are slightly more problematic, since their distribution is not isotropic, but they are relevant mostly at small scales, close to the central engine. There is also synchrotron- self-Compton emission from the radio-synchrotron photons themselves. For the first (CMB) and last (SSC) cases we can assume an isotropic distribution of photons (Hard- castle et al. 1998b). The emissivity equations are not derived in detail here, and can be found inRybicki & Lightman(1986).

While in FR II lobes inverse-Compton is the main contribution to the total energy budget, thermal processes contribute to the total energy as well. Shocks around FR II lobes have proved quite elusive, but there are some examples (see e.g. 3C 444Croston

et al. 2011). The thermal energy contribution for low power sources is much larger, and is studied in detail for two objects in this thesis (Chapters3and4). This energy can be derived from assumptions on the thermal spectrum and the shock conditions. Details on the interactions of the lobes with the environment, and shock physics, are given in Section1.2.